Datasets:
charlieoneill
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[ "Satyapal, Shobita", "Ellison, Sara L.", "McAlpine, William", "Hickox, Ryan C.", "Patton, David R.", "Mendel, J. Trevor" ]
2014MNRAS.441.1297S
[ "Galaxy pairs in the Sloan Digital Sky Survey - IX. Merger-induced AGN activity as traced by the Wide-field Infrared Survey Explorer" ]
196
[ "Department of Physics, Astronomy, & Computational Sciences, George Mason University, MS 3F3, 4400 University Drive, Fairfax, VA 22030, USA", "Department of Physics & Astronomy, University of Victoria, Finnerty Road, Victoria, BC V8P 1A1, Canada", "Department of Physics, Astronomy, & Computational Sciences, George Mason University, MS 3F3, 4400 University Drive, Fairfax, VA 22030, USA", "Department of Physics and Astronomy, Dartmouth College, 6127 Wilder Laboratory, Hanover, NH 03755, USA", "Department of Physics & Astronomy, Trent University, 1600 West Bank Drive, Peterborough, ON K9J 7B8, Canada", "Max-Planck-Institut fur Extraterrestrische Physik, Giessenbachstrasse, D-85748 Garching, Germany" ]
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[ "astronomy" ]
20
[ "galaxies: active", "galaxies: evolution", "Galaxies: interactions", "galaxies: Seyfert", "Astrophysics - Astrophysics of Galaxies", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
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[ "10.1093/mnras/stu650", "10.48550/arXiv.1403.7531" ]
1403
1403.7531_arXiv.txt
Based on the current cold dark matter cosmological framework, it is now well-established that galaxy interactions are ubiquitous and that they play a pivotal role in the formation and evolution of galaxies. From both a theoretical and observational perspective, galaxy interactions are undoubtedly responsible for enhanced nuclear star formation \citep[e.g.,][]{mihos1996,larson1978,sanders1996,kennicutt1987,per06,woods2006,woods2007,dim07,dim08,ellison2008,cox08,smith2010,patton2011,liu12,scudder2012,patton2013}, and the formation of spheroids \citep[e.g.,][]{toomre1977,lake1986,shier1998,rothberg2006}. A natural assumption from the tight correlation between central black hole mass and bulge velocity dispersion \citep[e.g.,][]{gebhardt2000} is that in addition to bulge growth, interactions trigger accretion onto a central supermassive black hole. However, despite over three decades of extensive research, it is still a topic of debate whether or not there is observational evidence for a causal connection between mergers and Active Galactic Nuclei (AGN), and, if so, how this connection depends on merger and host galaxy parameters. A number of studies have found evidence for mergers in luminous quasar hosts \citep[e.g.,][]{canalizo2001,benn08,urru08,for09,ramos2011,bess12,urru12}, where the link to interactions is fairly well accepted. On the other hand, the connection to the less luminous AGN population remains controversial. In particular, studies that have looked for an excess of tidal features or distorted morphologies in AGN versus non-AGN galaxies, have found no statistical difference both at low \citep[e.g.,][]{gab09,rei09,cisternas2011,koc11,bohm12}, intermediate(0.5$<$z$<$0.8) \citep[e.g.,][]{villforth2014} and higher (z$>$1) redshifts \citep[e.g.,][]{karouzos2014,fan2014} . Conversely, studies of close pairs have found enhanced fractions of AGN (or accretion rates), which supports a link between mergers and nuclear activity \citep[e.g.,][]{alo07,woods2007,koss2010,ellison2011,sil11,koss2012,liu12,ellison2013a,sbaf13}. This discrepancy may be due in part to the low surface brightness of tidal features and the time during the interaction at which they are expected to be visible \citep[e.g.][]{lotz08}. If the luminosity of the AGN is variable over a wide dynamic range on timescales shorter than the lifetime of merger signatures, any observed trends of merger fraction as a function of AGN luminosity will be weak, while the incidence of AGNs in merging galaxies will be still higher than in isolated systems \citep{hickox2014}. Furthermore, since tidal features can be faint and appear only in the gas instead of the stars \citep[e.g.][]{kuo2008}, the sensitivity \citep[e.g.][]{canalizo2001,ramos2011} and the wavelength \citep[e.g.,][]{hancock2007,boselli2005} of the observations may play a role in identifying merger signatures in AGN hosts. However, an alternative way of reconciling the apparently conflicting results is if mergers \textit{can} trigger AGN, but the majority of AGN are not produced through an interaction \citep[e.g.][]{db12}. It is now well-known that observations in only one waveband cannot provide a complete census of AGNs in galaxy samples due to obscuration of the central source or contamination of the observed emission by the host galaxy \citep[e.g.][]{satyapal2008,goulding2009,hickox2009,donley2010,juneau2013}. Although AGN excesses in samples of galaxy pairs have been found at optical, radio and X-ray wavelengths \citep[e.g.,][]{woods2007,koss2010,sil11,koss2012,liu12,sbaf13}, a direct comparison of these selection techniques has not been previously performed, and we have little understanding of what the complete census of merger-induced AGN might be. In our previous work on galaxy pairs and post-mergers in the Sloan Digital Sky Survey, we have used optical emission line diagnostics to identify an enhanced AGN fraction relative to a control that increases with decreasing pair separations \citep{ellison2011} and peaks post-coalescence \citep{ellison2013a}. However, since the centres of interacting galaxies may be more obscured than isolated galaxies, a potential limitation of this and previous optical studies, is that obscured AGNs can be missed at various stages along the merger sequence \citep{goulding2012}. In such cases, mid-infrared observations are a powerful tool for finding optically obscured AGNs. While there have been a number of mid-infrared studies of interacting galaxies, virtually all past studies have employed small samples of galaxies and/or have targeted the most advanced stage mergers \citep[e.g.,][]{genzel98,veilleux2009,armus07,armus09,farrah07,petric11}. The all-sky survey carried out by the {\it Wide-field Infrared Survey Explorer (WISE)} \citep{wright2010} has opened up a new window in the search for optically hidden AGNs in a large number of galaxies. This is because hot dust surrounding AGNs produces a strong mid-infrared continuum and infrared spectral energy distribution (SED) that is clearly distinguishable from star forming galaxies in both obscured and unobscured AGNs \citep[e.g.][] {lacy2004,stern2005,donley2007,stern2012}. The \textit{WISE} survey enables a more statistically significant study of the optically obscured AGN population in interacting galaxies. The goal of this paper is to complement our previous optical AGN study of SDSS galaxy pairs with a measurement of the incidence of obscured AGN, using mid-infrared colour selection with \textit{WISE}. This is the first large mid-infrared study of galaxy pairs. In Section \ref{sample_sec} we describe the selection of our samples of galaxy pairs, post-mergers, and their matched controls. In Section \ref{WISE_sec}, we discuss our \textit{WISE} AGN classification criteria, followed in Section \ref{IRAC_sec} by a discussion of the fidelity of our {\it WISE} photometry for close pairs. in Section \ref{results_sec}, we determine the mid-infrared colour-selected AGN fraction in the merger samples using \textit{WISE} compared to the control sample. In Section \ref{other_sec} we discuss other causes of red \textit{WISE} colours, followed by a summary of our results in Section \ref{summary_sec}. Throughout the manuscript we adopt a cosmology with $H_0$ = 70 \kms $,\Omega_M=0.3$, and $\Omega_{\Lambda}$=0.7.
\label{summary_sec} We have conducted a mid-infrared study aimed at finding obscured AGNs using {\it WISE} matched to a large sample of galaxy pairs and post-mergers selected from the SDSS. This is the first mid-infrared investigation of a large sample of galaxy pairs. Our main results can be summarized as follows: \begin{enumerate} \item{ We find a higher fraction of AGNs in galaxy pairs compared to a carefully constructed control sample of isolated galaxies matched in redshift, mass and local environment for pair separations less than 50 \hkpc\ .} \item{ The excess in the AGN fraction over the matched control increases with decreasing pair separation. The excess is most dramatic in the post-merger sample, where we find a factor of 10-20 excess in the AGN fraction compared to the control, depending on the adopted colour threshold.} \item{The trend of increasing infrared selected AGN fraction with decreasing pair separation, peaking in the post-merger sample, is in qualitative agreement with results based on AGN selection obtained from optical emission line diagnostics. However, the AGN excess based on infrared colour selection is significantly larger than the AGN excess based on optical spectroscopic diagnostics for the galaxy pairs at the smallest pair separations, with the most dramatic discrepancy seen in the post mergers. Our results imply that AGNs are both more energetically dominant and obscured with decreasing pair separation, as expected based on theoretical predictions. } \item{The AGNs in galaxy pairs show a significant enhancement of W2 band luminosity compared to their matched control out to at least 80 \hkpc, and is largest (by a factor of 3) for the post-mergers. This is consistent with the results from \citep{ellison2013a} which show a similar trend using the [OIII] luminosity of optical AGN indicative of enhanced black hole accretion rate for close pairs and post-mergers over isolated galaxies.} \end{enumerate}
14
3
1403.7531
Interactions between galaxies are predicted to cause gas inflows that can potentially trigger nuclear activity. Since the inflowing material can obscure the central regions of interacting galaxies, a potential limitation of previous optical studies is that obscured active galactic nuclei (AGNs) can be missed at various stages along the merger sequence. We present the first large mid-infrared study of AGNs in mergers and galaxy pairs, in order to quantify the incidence of obscured AGNs triggered by interactions. The sample consists of galaxy pairs and post-mergers drawn from the Sloan Digital Sky Survey that are matched to detections by the Wide-Field Infrared Sky Explorer. We find that the fraction of AGNs in the pairs, relative to a mass-, redshift- and environment-matched control sample, increases as a function of decreasing projected separation. This enhancement is most dramatic in the post-merger sample, where we find a factor of 10-20 excess in the AGN fraction compared with the control. Although this trend is in qualitative agreement with results based on optical AGN selection, the mid-infrared-selected AGN excess increases much more dramatically in the post-mergers than is seen for an optical AGN. Our results suggest that energetically dominant optically obscured AGNs become more prevalent in the most advanced mergers, consistent with theoretical predictions.
false
[ "optical AGN selection", "obscured AGNs", "mergers", "previous optical studies", "galaxy pairs", "projected separation", "nuclear activity", "AGN", "theoretical predictions", "interacting galaxies", "AGNs", "galaxies", "mergers and galaxy pairs", "active galactic nuclei", "various stages", "gas inflows", "the Sloan Digital Sky Survey", "-", "detections", "post" ]
14.72042
7.104412
119
792099
[ "Baldi, Marco" ]
2014PDU.....3....4B
[ "Cold dark matter halos in Multi-coupled Dark Energy cosmologies: Structural and statistical properties" ]
9
[ "-" ]
[ "2014PhRvD..90h3508A", "2016JCAP...01..050G", "2016Symm....8...70A", "2017JCAP...11..001M", "2017MNRAS.465..653B", "2018LRR....21....2A", "2018MNRAS.481.2933H", "2020PDU....3000666D", "2021EPJC...81..544A" ]
[ "astronomy" ]
7
[ "Astrophysics - Cosmology and Extragalactic Astrophysics" ]
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[ "10.1016/j.dark.2014.03.001", "10.48550/arXiv.1403.2408" ]
1403
1403.2408_arXiv.txt
\label{sec:intro} Any extension of the standard cosmological model that aims to overcome the fine-tuning problems of the cosmological constant must introduce additional degrees of freedom playing the role of Dark Energy (DE) in the form of a new dynamical field with defined clustering and interaction properties \citep[][]{Euclid_TWG} or as a low-energy modification of the laws of gravity. Despite the ever increasing accuracy of observational data of different kinds \citep[see e.g.][]{Vikhlinin_etal_2009b,Reid_etal_2010,Dunkley_etal_2010,wmap9,Planck_016} showing full consistency with the expected behaviour of a $\Lambda $CDM cosmology, such extended scenarios retain their appeal as the only possible alternative to anthropic arguments in explaining the observed value of the Dark Energy density at present. Nonetheless, large deviations from the standard background and linear perturbations evolution have been progressively ruled out by data, thereby significantly restricting the parameter space of a wide range of extended cosmologies to the level where potentially observable features would be ever hardly detectable. In this respect, a considerable interest has been attracted in the last years by extended cosmologies characterised by various types of screening mechanisms capable of recovering the standard $\Lambda $CDM behaviour in the appropriate regimes that are tightly constrained by observations, still allowing for possibly detectable deviations elsewhere. For modified theories of gravity (including e.g. $f(R)$ models) the recovery of the standard behaviour can be enforced both at cosmological scales, by selecting the model's parameters in order to reproduce as close as possible the standard $\Lambda $CDM expansion history \citep[see e.g.][]{Hu_Sawicki_2007}, and within overdense regions of the universe (like our Galaxy), through various types of screening mechanisms \citep[as the {\em chameleon}, {\em symmetron}, and {\em Vainshtein} mechanisms, see e.g.][respectively]{Khoury_Weltman_2004,Hinterbichler_Khoury_2010,Deffayet_etal_2002}. Such construction allows compatibility with both cosmological observations and solar-system tests of gravity \citep[see e.g.][]{Bertotti_Iess_Tortora_2003,Will_2005,Amendola_Tsujikawa_2008,Capozziello_Tsujikawa_2008}. Alternatively, if the new degree of freedom associated with DE has a selective interaction to Dark Matter, leaving baryonic particles uncoupled \citep[][]{Damour_Gibbons_Gundlach_1990}, as for the case of coupled Dark Energy (cDE) models \citep[][]{Wetterich_1995,Amendola_2000,Farrar2004}, solar system tests of gravity do not apply, and the model is mainly constrained by its impact on cosmological observables \citep[][]{Bean_etal_2008,Xia_2009,Baldi_Viel_2010,Pettorino_etal_2012}. However, even in this case such constraints are capable to bind the DE-CDM interaction to a level that makes its predicted observational signatures at small scales hardly detectable with the presently available observational precision, at least for the most widely considered case of a constant coupling. Several simple extensions of such standard cDE scenario have been proposed in recent years, with the aim to explore the range of plausible non-standard DE cosmologies that might allow for testable predictions at the scales of structure formation without conflicting with the ever more stringent bounds arising from cosmological probes. These extensions include the possibility of time-dependent coupling functions \citep[see e.g.][]{Amendola_2004,Baldi_2011a}, resulting in a strongly suppressed impact of the interaction on the background cosmic evolution even in the presence of significant effects on the formation and evolution of nonlinear structures; alternatively, the possibility of having multiple species of CDM particles interacting with DE through individual coupling functions have also been proposed \citep[see e.g.][]{Brookfield_VanDeBruck_Hall_2008}. The latter scenarios are characterised by attractor solutions that suppress the effective coupling to DE during matter domination, thereby making the model consistent with the observed expansion history \citep[][]{Piloyan_etal_2013}. Although such ``extended" cDE models might appear in general less appealing than standard quintessence and cDE scenarios -- as they generally involve one or more additional free parameters associated with the specific time evolution of the coupling function or to the individual couplings of different CDM species -- a particular realisation of the latter class of models requiring no additional parameters as compared to standard cDE cosmologies has been recently proposed \citep[][]{Baldi_2012a}, and termed the ``Multi-coupled DE" (McDE, hereafter) scenario. This features only two distinct CDM particle species with opposite constant couplings to a DE scalar field, and has been shown to provide a very effective screening of the interaction during matter domination even for very large values of the coupling constant $\beta $ \citep[see again][]{Piloyan_etal_2013}. Such screening is then broken at the onset of DE domination, thereby providing a time-dependent effective coupling without imposing {\em a priori} any time evolution of the coupling function. {It is important to stress again here that the idea of a non-universal coupling between a cosmic scalar and distinct matter fluids is not a new concept that characterises only the McDE scenario: indeed, such idea is at the very basis of the whole class of Coupled Quintessence models, where a non-universal coupling has to be invoked \citep[see again][]{Damour_Gibbons_Gundlach_1990} in order to keep baryonic particles minimally-coupled so to evade constraints on scalar-tensor theories from local tests of gravity. The truly distinctive feature of the McDE scenario, however,} is that the presence of {coupling constants with an opposite sign} for the two different CDM particle species determines the existence of both attractive and {\em repulsive} fifth-forces between CDM particles, differently from all standard cDE models where fifth-forces are always attractive. {It is also worth mentioning how the possibility of a multi-particle nature of the CDM cosmic fraction has been proposed in several different theoretical contexts \citep[see e.g. the recent work by][]{Chialva_Dev_Mazumdar_2013}, considering matter species that differ from each other in various physical properties. For instance, the possibility of multi-particle dark matter models featuring a cold and a hot component has been proposed \citep[see e.g.][]{Anderhalden_etal_2012,Maccio_etal_2013} as a possible solution of the small-scale problems of the $\Lambda $CDM scenario, which in some specific realisations include also a non-vanishing coupling of the cold species to a DE scalar field \citep[as for the model proposed by][]{Bonometto_Sassi_LaVacca_2012,Bonometto_Mainini_2014}. At a more fundamental level, models with multiple CDM species with different couplings to a cosmic light scalar might arise in the context of string-inspired cosmological scenarios \citep[see e.g.][]{Brandenberger_Vafa_1989} as proposed by \cite{Gubser_Peebles_2004} and subsequently extensively discussed by \cite{Gubser2004,Farrar2004,Nusser_Gubser_Peebles_2005}. The McDE scenario, in particular, represents the simplest realisation of the model proposed by \cite{Gubser_Peebles_2004} which has as its most characteristic footprint the existence of long-range attractive and repulsive fifth-forces.} {For coupling values that appear to be consistent with the observed background and linear perturbations evolution \citep[see e.g. the recent work by][]{Piloyan_etal_2014}, such long-range fifth-forces might have a strength comparable to standard gravity, giving rise to a very peculiar phenomenology at the level of linear and nonlinear structure formation \citep[][]{Baldi_2013}. The work by \cite{Piloyan_etal_2014}, however, has also shown how the evolution of linear density perturbations suddenly deviates from the standard $\Lambda $CDM behaviour when the coupling grows beyond the value corresponding to a gravitational strength of the associated fifth-forces, thereby allowing to place much tighter constraints on the coupling itself through measurements of the linear growth rate as compared to the extremely loose bounds derived from the background expansion history alone.} This phenomenology has been already explored with both analytical and numerical tools, and in particular the nonlinear evolution of large-scale structures has been investigated with some first low-resolution N-body simulations \citep[][]{Baldi_2013} aimed at coarsely sampling the model's parameter space and highlighting its most prominent effects on the large-scale shape of the cosmic web. Such analysis has shown that fifth-forces with the same strength as standard gravity appear to still have a relatively mild impact on the overall distribution of nonlinear structures in McDE cosmologies, and allowed to observe for the first time the halo fragmentation process associated with the repulsive interaction between the two different CDM particle species. The details of such small scale effects, however, could not be observed due to the limited resolution of those early N-body simulations, and their proper investigation demands higher resolution runs for the range of parameters that already showed to provide a reasonable behaviour of structure formation at large scales. With the present work, we move in such direction by presenting the first high-resolution N-body simulations of McDE models ever performed, and investigating the effects of such cosmologies on the statistical and structural properties of CDM halos for several different values of the coupling constant $\beta $. \\ The paper is organised as follows. In Section~\ref{sec:McDE} we review the main features of McDE models at the background and linear perturbations level; in Section~\ref{sec:sims} we describe the numerical setup of our high-resolution N-body simulations, and illustrate the basic post-processing analysis performed on the simulations outputs; in Section~\ref{sec:results} we discuss the results of our investigation on a number of potentially observable statistical and structural properties of CDM halos. Finally, in Section~\ref{sec:concl} we drive our conclusions.
\label{sec:concl} In this paper, we have presented the outcomes of the first high-resolution N-body simulations of Multi-coupeld Dark Energy cosmologies. These are cosmological models characterised by the existence of two distinct species of CDM particles with opposite couplings to a classical Dark Energy scalar field, giving rise to both {\em attractive} and {\em repulsive} long-range scalar interactions between CDM particles. While requiring the same number of parameters as a standard coupled quintessence model, Multi-coupled Dark Energy scenarios provide a very effective way to screen the coupling during matter domination, thereby strongly alleviating the impact of the interaction between Dark Energy and CDM particles on the background and linear perturbations evolution of the universe. In this respect, Multi-coupled Dark Energy models are practically indistinguishable from a standard $\Lambda $CDM cosmology up to very recent epochs for a wide range of couplings. Nonetheless, the effects of the additional fifth-forces in the nonlinear regime of structure formation are expected to imprint characteristic features in the statistical and structural properties of CDM halos that might allow to observationally test the model. To this end, a first series of low-resolution N-body simulations was performed by \citet{Baldi_2012b} with the main purpose of sampling the parameter space of the model and highlight such possible observational footprints. In this work, we did improve with respect to those early simulations in several aspects, by running the first high-resolution simulations of Multi-coupled Dark Energy models that allow to investigate in detail the statistical and structural properties of CDM halos arising in these cosmologies up to a coupling value of $|\beta |=1$. We briefly recap here the main conclusions of our study.\\ First of all, the various effects that we have highlighted with our high-resolution simulations, and that will be listed below, show a very strong dependence on the coupling value, and in particular on whether the coupling is below or above the gravitational coupling $|\beta |=\sqrt{3}/2$. Such value defines the threshold between the regimes where scalar fifth-forces are weaker ($|\beta |< \sqrt{3}/2$) or stronger ($|\beta |> \sqrt{3}/2$) than standard gravity. In general, for all the coupling values of our sample lying in the former range we found very mild effects of the DE-CDM interaction on all the statistical and structural properties of CDM halos that we have investigated; on the contrary, the impact of the interaction becomes very quickly extremely significant as soon as the gravitational coupling threshold is reached and overcome. Therefore, the gravitational coupling $|\beta |=\sqrt{3}/2$ represents an intrinsic threshold for Multi-coupled Dark Energy models, and the small range of coupling values around such threshold promise to determine the most interesting phenomenological effects of such scenarios, providing at the same time viable and non-trivial observational effects: coupling values much below this threshold appear to be almost indistinguishable from a standard $\Lambda $CDM cosmology, while values significantly above it determine very dramatic effects at nonlinear scales that are likely to be easily ruled out by presently available data. Such effects, according to our present analysis, can be summarised as follows. \begin{center} {\em -- Large-scale density smoothing and halo fragmentation -- } \end{center} At the largest scales included in our cosmological simulations ($100$ Mpc$/h$ aside) the main effect of the interaction between Dark Energy and CDM in Multi-coupled Dark Energy cosmologies is to smooth the density field of the total CDM fluid, consistently with the earler findings \citep[][]{Baldi_2013} of a significant suppression of power at mildly nonlinear and nonlinear scales. Such effect, however, becomes appreciable only for the largest coupling value included in our sample, $|\beta |=1$, and is almost absent for all the other coupling values considered in our work. At smaller scales, the most prominent effect of the interaction is the fragmentation of bound CDM structures into smaller objects as a consequence of the different dynamical evolution characterising the two distinct CDM particle species. As anticipated above, such phenomenon is completely absent for coupling values below the gravitational coupling threshold $|\beta | < \sqrt{3}/2$, as the gravitational attraction between CDM particles of opposite type is still stronger than their fifth-force repulsion, and sufficient to overcome the effect of the scalar friction that would tend to drag the two different CDM species in opposite directions along their unperturbed trajectory. However, already for a coupling $|\beta |=\sqrt{3}/2$, our simulations have shown the fragmentation of individual structures into pairs of objects of comparable size, and how the separation of such fragmented halos grows in time along the trajectory of the original parent structure, thereby providing a direct evidence of the violation of the weak equivalence principle in Multi-coupled Dark Energy cosmologies. It is particularly relevant to stress here that even the largest coupling value below the gravitational coupling threshold in our set of models, $|\beta |=8/10$, did not show any sign of halo fragmentation. This gives a feeling of how sharp the transition between the two regimes around the gravitational coupling threshold is. \begin{center} {\em -- Halo mass function and suppression of cluster abundance --} \end{center} We have computed the abundance of CDM halos as a function of their mass in all the cosmological models under investigation and at different redshifts, based on a halo catalogue computed through a Friends-of-Friends algorithm without any distinction between the two CDM particle species. By comparing the obtained halo mass functions to the uncoupled case, we could highlight a series of characteristic footprints of Multi-coupled Dark Energy scenarios. Even in this case, such features are either completely absent or extremely weak for coupling values below the gravitational coupling threshold, while becoming very significant at and above the threshold. First of all, as a consequence of the halo fragmentation process there is a clear reduction of the abundance of intermediate mass halos, and a corresponding enhancement of smaller mass halos. Relative to the uncoupled case, the latter effect has roughly twice the amplitude of the former, since the fragmentation process occurs by splitting an originally mixed CDM object into two separate halos dominated by the two different CDM particle types, respectively, and with roughly half the mass of the parent structure, such that for any intermediate mass halo that disappears two new half-mass objects will populate the low-mass end of the halo mass function. Our results also showed for the first time how the halo fragmentation process evolves in a hierarchical fashion, with halos of small and intermediate mass starting to fragment first, followed by progressively larger mass halos at lower redshifts. This is clearly shown by the shift of the transition between low-mass enhancement and high-mass suppression of the halo abundance to progressively larger masses for decreasing redshifts. Clearly, the strong enhancement of low-mass CDM halos in Multi-coupled Dark Energy cosmologies might be inconsistent with the observed abundance of satellite galaxies at small scales, thereby allowing to put constraints on the model. It is however important to point out here that it is not obvious how the baryonic and stellar components might evolve during the fragmentation process of their host CDM halo, thereby making it difficult to naively apply the standard abundance matching between visible galaxies and CDM halos also in the context of Multi-coupeld Dark Energy models. Dedicated radiative hydrodynamical simulations with cooling and star formation would be required to assess this point, and will be pursued in future works. Another very interesting effect of Multi-coupled Dark Energy models on the abundance of collapsed structures concerns the mass range of galaxy clusters for large values of the coupling. In particular, for a coupling of order unity our simulations have shown a very significant suppression of the abundance of cluster sized CDM halos at low redshifts. Such feature might alleviate the present tension between the cosmologically inferred value of $\sigma _{8}$ and its best-fit value based on cluster counts, as reported also by the recent results of the Planck satellite mission. Clearly, a coupling value of $|\beta | = 1$ for Multi-coupled Dark Energy models might be disfavoured by other observable features such as e.g. the strongly enhanced abundance of low-mass halos. However, it is interesting to notice that our cosmological scenario might determine a suppression of the abundance of massive clusters without changing the large-scale normalisation of linear density perturbations, as such effect cannot be obtained in most of the other available Dark Energy or Modified Gravity models. \begin{center} {\em -- Structural properties of CDM halos --} \end{center} By matching individual structures in the different simulations, we have compared the structural properties of halos at different masses, taking the spherically averaged mass distribution around the most bound particle of each halo. Even in this case, our results have shown how for coupling values below the gravitational coupling threshold the overall impact of the DE-CDM interaction on the structural properties of halos is very mild, while it becomes significant for larger values of the coupling. In particular, we have investigated the total CDM density profile of halos, showing that for sub-gravitational couplings the density profile is almost unaffected and consistent with a standard NFW universal shape as for an uncoupled cosmology. On the contrary, larger values of the coupling determine a significant suppression of the overdensity in the inner part of the halos, and a corresponding increase of the density in the outer regions. Such effect is related to the outflow of the sub-dominant CDM species during the halo fragmentation process. We also found that, consistently with the previous results on the onset of the halo fragmentation process, halos in sub-gravitational coupling models are still composed by a mixture of the two distinct CDM particle species, while above the gravitational coupling threshold halos are dominated by a single CDM type at $z=0$. The resulting total gravitational potential (i.e. the potential arising by the superposition of the standard Newtonian potential and of the attractive and repulsive potentials associated to the scalar fifth-forces) has a non-trivial shape around CDM halos which depends on the relative distribution of the two CDM species within each structure. This effect gives rise to a species-dependent and (consequently) space-dependent screening mechanism of the scalar fifth-force even at the nonlinear and highly nonlinear level. Differently from other screening mechanisms associated to various classes of modified gravity theories, however, the screening mechanism of Multi-coupled Dark Energy cosmologies is a transient phenomenon, as it is based on the balance between attractive and repulsive corrections to standard gravity, and is therefore no longer effective as soon as mixed-CDM halos fragment into single-species objects. Nonetheless, for sub-gravitational coupling models such screening mechanism can hold until the present time, and provides a very effective suppression of the scalar fifth-forces.\\ To conclude, we have presented in this paper the first high-resolution N-body simulations of the Multi-coupled Dark Energy scenario, and investigated how such cosmological models affect the statistical and structural properties of collapsed objects forming from primordial density perturbations. Most remarkably, we have shown how the formation and evolution of cosmic structures is practically indistinguishable from a $\Lambda $CDM cosmology for coupling values below the gravitational coupling threshold $|\beta |=\sqrt{3}/2$. When such barrier is reached and overcome, our results have shown a series of very significant effects on structure formation, ranging from the fragmentation of previously bound halos into smaller objects, to the explicit violation of the weak equivalence principle, to the distortion of the universal shape of the halo mass function, to the modification of the density and gravitational potential profiles of halos that tend to become less overdense in their core. Finally, we have shown how Multi-coupled Dark Energy models feature a new type of nonlinear screening mechanism of the scalar fifth-forces associated with the coupling to Dark Energy, and how such screening mechanism is an unstable transient phenomenon that breaks down at the onset of halo fragmentation.
14
3
1403.2408
The recently proposed Multi-coupled Dark Energy (McDE) scenario - characterised by two distinct cold dark matter (CDM) particle species with opposite couplings to a Dark Energy scalar field - introduces a number of novel features in the small-scale dynamics of cosmic structures, most noticeably the simultaneous existence of both attractive and repulsive fifth-forces. Such small-scale features are expected to imprint possibly observable footprints on nonlinear cosmic structures, that might provide a direct way to test the scenario. In order to unveil such footprints, we have performed the first suite of high-resolution N-body simulations of McDE cosmologies, covering the coupling range |β| ≤ 1. We find that for coupling values corresponding to fifth-forces weaker than standard gravity, the impact on structure formation is very mild, thereby showing a new type of screening mechanism for long-range scalar interactions. On the contrary, for fifth-forces comparable to or stronger than standard gravity a number of effects appear in the statistical and structural properties of CDM halos. Collapsed structures start to fragment into pairs of smaller objects that move on different trajectories, providing a direct evidence of the violation of the weak equivalence principle. Consequently, the relative abundance of halos of different masses is significantly modified. For sufficiently large coupling values, the expected number of clusters is strongly suppressed, which might alleviate the present tension between CMB- and cluster-based cosmological constraints. Finally, the internal structure of halos is also modified, with a significant suppression of the inner overdensity, and a progressive segregation of the two CDM species.
false
[ "CDM halos", "opposite couplings", "Dark Energy", "CDM", "nonlinear cosmic structures", "cosmic structures", "|β| ≤", "standard gravity", "structure formation", "novel features", "Collapsed structures", "halos", "smaller objects", "different trajectories", "screening mechanism", "different masses", "McDE cosmologies", "long-range scalar interactions", "two distinct cold dark matter", "clusters" ]
11.243656
0.927585
89
513694
[ "Scafetta, Nicola", "Willson, Richard C." ]
2014Ap&SS.350..421S
[ "ACRIM total solar irradiance satellite composite validation versus TSI proxy models" ]
38
[ "Active Cavity Radiometer Irradiance Monitor (ACRIM), Coronado, CA, USA; Duke University, Durham, NC, USA", "Active Cavity Radiometer Irradiance Monitor (ACRIM), Coronado, CA, USA" ]
[ "2014Ap&SS.352..341W", "2014Ap&SS.354..275S", "2014PhyA..413..329S", "2015ESRv..150..409S", "2015JASTP.133..111R", "2015LRSP...12....4H", "2015NewA...34..221V", "2016AdSpR..58.2104V", "2016IAUS..320..342R", "2017JSWSC...7A..11W", "2017NewA...51..135Y", "2017PhyA..468..351K", "2018NewA...62...62R", "2018SoEn..168....2G", "2019AdSpR..64..104V", "2019ESRv..19802950S", "2019JApA...40...11Z", "2019RemS...11.2569S", "2020E&SS....700645L", "2020E3SWC.15405003G", "2020Ge&Ae..59..908V", "2020GeoRL..4789954S", "2020SoPh..295...20B", "2021Atmos..12..147S", "2021Atmos..12.1145S", "2021Clim....9..163S", "2021IJAA...11..279S", "2021RAA....21..131C", "2022Atmos..13.1132V", "2022FrP....10.9738C", "2022RAA....22l5008R", "2023Atmos..14.1339V", "2023Clim...11...77S", "2023Clim...11..179S", "2023GeoFr..1401650S", "2023JGRD..12838941A", "2024JHEP...02..064G", "2024SoPh..299...67D" ]
[ "astronomy", "physics", "earth science" ]
12
[ "Solar luminosity", "Total solar irradiance (TSI)", "Satellite experimental measurements", "TSI satellite composites", "TSI proxy model comparisons", "Physics - Geophysics", "Astrophysics - Solar and Stellar Astrophysics" ]
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[ "10.1007/s10509-013-1775-9", "10.48550/arXiv.1403.7194" ]
1403
1403.7194_arXiv.txt
The satellite total solar irradiance (TSI) database is now more than three and a half decades long and provides a valuable record for investigating the relative significance of natural and anthropogenic forcing of climate change \citep{IPCC,Scafetta2009,Scafetta2011}. It is made of 7 major independent measurements covering different periods since 1978 (see Figure 1). \begin{figure*}[!t] \centering \includegraphics[width=1.8\columnwidth]{figure1} \caption{Total solar irradiance satellite record database. } \end{figure*} A composite TSI record can be constructed from the series of experiments since 1978 by combining and cross-calibrating the set of overlapping satellite observations to create a TSI time series. TSI satellite composites provide end-to-end traceability at the mutual precision level of the overlapping satellite experiments that is orders of magnitude smaller than the absolute uncertainty of the individual experiments. The scale offsets of the various satellite results shown in Figure 1 are caused by the uncertainties of their self-calibration \citep{Willson2003,Frohlich2012}. Different approaches in selecting results and cross-calibrating the satellite records on a common scale have resulted in composites with different characteristics. \begin{figure*}[!t] \centering \includegraphics[width=1.8\columnwidth]{figure2}\caption{TSI satellite composites ACRIM and PMOD. Different approaches to bridging the ACRIM Gap result in different trends. The ACRIM composite uses: (1) ERB, ACRIM1, 2 and 3 results published by the experiment science teams; (2) ERB comparisons to bridge the ACRIM Gap; (3) ACRIM3 scale. The PMOD composite uses: ERB, ACRIM1, ACRIM2 and VIRGO results; (2) ERBE comparisons to bridge the ACRIM Gap; (3) Alters published ERB and ACRIM1 results to conform them to TSI proxy models; (4) VIRGO scale. } \end{figure*} Figure 2 shows the two TSI satellite composites most commonly cited: ACRIM \citep{Willson1997,Willson2001,Willson2003} and PMOD \citep{Frohlich1998,Frohlich2004,Frohlich2006,Frohlich2012}. Alternative TSI satellite composites have been proposed by \citet{Dewitte} and \citet{Scafetta2011} using different methodologies to merge the datasets. The new ACRIM composite uses the updated ACRI M3 record. ACRIM3 data was reprocessed after implementing corrections for scattering and diffraction found during recent testing and some other algorithm updates. The testing was performed at the TSI Radiation Facility (TRF) of the Laboratory for Atmospheric and Space Physics (LASP) \citep[\url{http://lasp.colorado.edu/home/}]{Kopp}. Two additional algorithm updates were implemented that more accurately account for instrument thermal behavior and parsing of shutter cycle data. These removed a component of the quasi-annual signal from the data and increased the signal to noise ratio of the data, respectively. The net effect of these corrections decreased the average ACRIM3 TSI value from $\sim1366$ $W/m^{2}$ \citep[see: ][]{Willson2003} to $\sim1361$ $W/m^{2}$ without affecting the trending in the ACRIM Composite TSI. Differences between ACRIM and PMOD TSI composites are evident, but the most obvious and significant one is the solar minimum-to-minimum trends during solar cycles 21 to 23. ACRIM presents a bi-decadal increase of +0.037\%/decade from 1980 to 2000 and a decrease thereafter. PMOD presents a steady multi-decadal decrease since 1978 (see Figure 2). Other significant differences can be seen during the peak of solar cycles 21 and 22. These arise from the fact that ACRIM uses the original TSI results published by the satellite experiment teams while PMOD significantly modifies some results to conform them to specific TSI proxy models \citep{Frohlich1998,Frohlich2004,Frohlich2006,Frohlich2012}. The single greatest challenge in constructing a precise composite extending before 1991 is providing continuity across the two-year ACRIM Gap (1989.53\textendash{} 1991.76) between the results of SMM/ACRIM1 \citep{Willson1991} and UARS/ACRIM2 \citep{Willson1994,Willson1997}. During this period the only observations available were those of the Nimbus7/ERB (hereafter referred to as ERB) \citep{Hoyt1992} and ERBS/ERBE (hereafter referred to as ERBE) \citep{Lee}. These experiments provided TSI observations that met the needs of the Earth Radiation Budget investigations at that time, but were less precise and accurate than the ACRIM experiments that were designed specifically to provide the long term precision and traceability required by climate and solar physics investigations. ACRIM1 and ACRIM2 were intended to overlap initiating an ACRIM TSI monitoring strategy designed to provide long term TSI traceability of results through the precision of on-orbit comparisons. ACRIM2 was delayed by the Challenger disaster, however, and eventually deployed two years after the last data from ACRIM1. This period is known as the ACRIM GAP (1989.5 - 1991.75), as shown in Figure 1. ACRIM1, ACRIM2 and ACRIM3 were dedicated TSI monitoring experiments capable of highly precise observations by virtue of their design and operation, which includes continuous electronic self-calibration, high duty cycle solar observations (ACRIM1: 55 min./orbit; ACRIM2: 35 min./orbit; ACRIM3: up to full sun during its 96 minute sun-synchronous orbit), sensor degradation self-calibration, high observational cadence ($~2$ minutes) and precise solar pointing. ERB and ERBE were less accurate and precise experiments designed to meet the less stringent data requirements of Earth Radiation Budget modeling. They were able to self-calibrate only infrequently (every 14 days), had limited solar observational opportunities (ERB: 5 min/orbit daily; ERBE: 5 minutes every 14 days, usually) and were not independently solar pointed, observing while the sun moved through their fields of view, all of which degraded their precision and accuracy. Bridging the ACRIM Gap using ERB and ERBE results is problematical not only because of their lower data quality but also because their results yield significantly different and incompatible trends during the ACRIM Gap. During the ACRIM Gap ERB results trend upward (linear regression slope = $0.27\pm0.04$ $Wm^{-2}/year$) while ERBE trend downward (linear regression slope =$-0.27\pm0.15$ $Wm^{-2}/year$). This causes the difference between the ACRIM and PMOD TSI trends during solar cycles 21-23. The ACRIM TSI composite uses unaltered ERB results to relate ACRIM1 and ACRIM2 records, while PMOD uses an altered ERB record based on some theoretical model predictions that better agree with the downward trend of the ERBE record during the ACRIM Gap. \begin{figure*}[!t] \centering \includegraphics[width=1.4\columnwidth, angle=-90]{figure3}\caption{{[}A{]} The \citet{Krivova2007} magnetic field proxy model. {[}B{]} \citet{Wenzler2006} (WSKF06) standard surface magnetic field proxy models. {[}C{]} \citet{Wenzler2006} (WSKF06) optimized (on PMOD) standard surface magnetic field proxy models. The 1979 peak maximum (red letter ``P'') in the WSKF06 models and the lack of available data in 1992 (red letter \textquotedblleft{}V\textquotedblright{}) that separates NSO-512 (Feb/1/74 to Apr/18/92 and Nov/28/92 to Apr/10/93) and NSO-SPM (Nov/21/92 to Sep/21/03) records calls into question the accuracy of their cross-calibration with the KBS07 model. } \end{figure*} In Section 2 we review the hypotheses proposed in the literature about Nimbus7/ERB TSI record during the ACRIM Gap. In sections 3-8 we test these hypotheses by directly comparing the Nimbus7/ERB data sets versus alternative solar data and proxy models. In this process we will study in details the TSI proxy models of \citet{Krivova2007} (KBS07), and \citet{Wenzler2006} (WSKF06) shown in Figure 3. In Appendix A Hoyt (the head of the NASA Nimbus7/ERB science team) explains the accuracy of the ERB record during the ACRIM Gap. Appendix B briefly summarizes the importance of the TSI satellite composite issue for solar physics and climate change.
We have conducted several independent evaluations of the accuracy of TSI satellite data and their composites. The ACRIM TSI composite relies solely on the continuity of the results of overlapping satellite experiments as understood and published by the flight experiment teams. The ACRIM composite has a direct and exclusively experimental justification \citep{Willson2003}. On the contrary, the PMOD TSI composite \citep{Frohlich1998,Frohlich2004,Frohlich2006,Frohlich2012} is essentially a theoretical model originally designed to agree with Lean\textquoteright{}s TSI proxy model \citep{Frohlich1998}. It relies on postulated but experimentally unverified drifts in the ERB record during the ACRIM Gap, and other alterations of the published ERB and ACRIM results, that are not recognized by their original experimental teams and have not been verified by the PMOD by original computations using ERB or ACRIM1 data. Our findings support the reliability of the ACRIM composite as the most likely and precise representation of 35 years of TSI monitoring by satellite experiments. The only caveat is that the ERB record prior to 1980 may require some correction for degradation, but it would be much less than used in the PMOD composite. We argued that the ACRIM composite most closely represents true TSI because the very corrections of the published TSI data made by Fr\"ohlich to construct the PMOD composite are not supported by a direct comparison between ERBE and ERB records in the proximity of September/October 1989. Direct comparison of ERB and ERBE during 1989 showed that Fr\"ohlich\textquoteright{}s postulated Sep/29/1989 step function increase of $0.4$ $W/m^{2}$ in ERB sensitivity, which coincided with a power down event, did not occur. The KBS07 proxy model does not support Fr\"ohlich\textquoteright{}s ERB \textquoteleft{}glitch\textquoteright{} either. A divergence between the two satellite records did occur in November 1989; but this is more than one month later and clearly not associated with the ERB end-of-September power down event. We have demonstrated that the update of Lean\textquoteright{}s TSI proxy model \citep{Kopp2011}, used originally to validate PMOD\textquoteright{}s lack of trending from 1980 to 2000 \citep{Frohlich1998}, has inadequate predictive capability to properly reconstruct the TSI decadal trending. Lean\textquoteright{}s model predicted an upward trend between the TSI minima in 1996 and 2008 while both ACRIM and PMOD present a downward trend. This demonstrates that Lean's proxy model cannot reconstruct TSI decadal trending with a precision smaller than $\pm0.5$ $W/m^{2}$ the same order as the difference between PMOD and ACRIM TSI composites. Thus, the use of Lean's TSI proxy model is not useful as a guide to correct satellite measurements. The WSKF06 TSI proxy models contradict the primary PMOD rationale by the following findings: (1) there was a TSI peak in late 1978 and early 1979 as recorded by ERB (although some early mission degradation of the instrument may have been uncompensated for); (2) the ACRIM1 published record is more stable than ERB during 1980-1984 and should be preferred for constructing a TSI composite during this period; (3) ERB did not experience either the end-September 1989 step function drift in sensitivity or the upward linear drift claimed by Fr\"ohlich during 1990-1992.5. The latter result is also evident in the upgraded SATIRE model \citep{Ball}. Thus, if ERB requires some correction during the ACRIM Gap, our results suggest that Fr\"ohlich overestimated those corrections by at least a factor of two due to the fact that at least one of the two hypotheses (the ERB glitch in Sep/29/1989 or the ERB drift from Oct/1989 to 1992) are not confirmed by our cross-analysis. The ERB-ERBE divergence during the ACRIM Gap most likely resulted from uncorrected degradation of ERBE in its first exposure to short wavelength fluxes driven by enhanced solar activity during the 1989-1993 solar maximum or other events. Consequently PMOD should be shifted upward by about $0.5$ $W/m^{2}$ after 1992 which produces a 1980-2000 TSI upward trending similar to that observed in the ACRIM composite. Our results demonstrated that the validity of TSI proxy models should not be overestimated since they frequently produce conflicting results and contradictory features. Although Solanki's TSI proxy models appear to reproduce the lack of a trend during solar cycles 21 - 22 in the PMOD TSI composite, they contradict one or more of the hypotheses advocated by PMOD to alter the originally published TSI used in constructing the PMOD TSI composite. Thus some of the arguments used to promote the PMOD composite \citep{Frohlich1998,Wenzler2006,Krivova2007,Wenzler2009} are little more than speculations and coincidences. Moreover, Lean\textquoteright{}s and Solanki\textquoteright{}s TSI models differ significantly from the TSI model proposed by \citet{Hoy1993} who constructed a TSI record since 1700 using five alternative solar irradiance proxy indexes \textemdash{} sunspot cycle amplitude, sunspot cycle length, solar equatorial rotation rate, fraction of penumbral spots, and the decay rate of the sun spot cycle. Although Lean\textquoteright{}s and Solanki\textquoteright{}s models present no TSI trend during 1980-2000, as shown in the PMOD composite, there are other studies suggesting generally increasing TSI from 1970 to 2000. \citet{Shapiro} found a small increasing trend across the 1975, 1986 and 1996 solar minima followed by a decrease in the minimum of 2008. Also the cosmic ray flux index would suggest a solar activity increase from 1980 to 1996 \citep[figure 20]{scafetta2013c}. The TSI pattern revealed in the ACRIM satellite composite is consistent with a quasi 60-year solar cycle modulation, which appears to be one of the major harmonic constituents of solar activity and should have theoretically peaked around 2000 \citep{Ogurtsov,Scafetta2012a,Scafetta2013a,Scafetta2013EE}. We conclude that solar activity may have presented a larger secular variability and specific geometrical patterns that are quite different from the Lean TSI model currently used to force the CMIP5 models. \subsection*{Acknowledgment: } The National Aeronautics and Space Administration supported Dr. Willson under contracts NNG004HZ42C at Columbia University and Subcontracts 1345042 and 1405003 at the Jet Propulsion Laboratory.
14
3
1403.7194
The satellite total solar irradiance (TSI) database provides a valuable record for investigating models of solar variation used to interpret climate changes. The 35-year ACRIM total solar irradiance (TSI) satellite composite time series has been revised using algorithm updates based on 13 years of accumulated mission experience and corrections to ACRIMSAT/ACRIM3 results for scattering and diffraction derived from recent testing at the Laboratory for Atmospheric and Space Physics/Total solar irradiance Radiometer Facility (LASP/TRF). The net correction lowers the ACRIM3 scale by ∼3000 ppm, in closer agreement with the scale of SORCE/TIM results (average total solar irradiance ≈1361.5 W/m<SUP>2</SUP>). Differences between the ACRIM and PMOD TSI composites are investigated, particularly the decadal trending during solar cycles 21-22 and the Nimbus7/ERB and ERBS/ERBE results available to bridge the ACRIM Gap (1989-1992), are tested against a set of solar proxy models. Our findings confirm the following ACRIM TSI composite features: (1) The validity of the TSI peak in the originally published ERB results in early 1979 during solar cycle 21; (2) The correctness of originally published ACRIM1 results during the SMM spin mode (1981-1984); (3) The upward trend of originally published ERB results during the ACRIM Gap; (4) The occurrence of a significant upward TSI trend between the minima of solar cycles 21 and 22 and (5) a decreasing trend during solar cycles 22-23. The same analytical approach does not support some important features of the PMOD TSI composite: (1) The downward corrections applied to the originally published ERB and ACRIM1 results during solar cycle 21; (2) The step function sensitivity change in ERB results at the end-of-September 1989; (3) The downward trend of ERBE results during the ACRIM Gap and (4) the use of ERBE results to bridge the ACRIM Gap. Our analysis provides a first order validation of the ACRIM TSI composite approach and its 0.037 %/decade upward trend during solar cycles 21-22. The implications of increasing TSI during the global warming of the last two decades of the 20th century are that solar forcing of climate change may be a significantly larger factor than represented in the CMIP5 general circulation climate models.
false
[ "solar cycle", "solar cycles", "solar proxy models", "solar variation", "average total solar irradiance", "ERB results", "solar forcing", "ERBE results", "climate change", "climate changes", "results", "upward trend", "TSI composites", "TSI", "models", "ERB", "originally published ERB results", "the CMIP5 general circulation climate models", "originally published ACRIM1 results", "a significant upward TSI trend" ]
11.17215
15.627987
2
482919
[ "Chen, Che-Yu", "Ostriker, Eve C." ]
2014ApJ...785...69C
[ "Formation of Magnetized Prestellar Cores with Ambipolar Diffusion and Turbulence" ]
104
[ "Department of Astronomy, University of Maryland, College Park, MD 20742, USA", "Department of Astronomy, University of Maryland, College Park, MD 20742, USA; Department of Astrophysical Sciences, Princeton University, Princeton, NJ 08544, USA;" ]
[ "2014ApJ...790L..19F", "2014ApJ...794..165S", "2014MNRAS.443..230H", "2014PhDT.......337H", "2015ApJ...801...77M", "2015ApJ...801..117T", "2015ApJ...804...37L", "2015ApJ...805....1B", "2015ApJ...805..185S", "2015ApJ...806...31G", "2015ApJ...806..226M", "2015ApJ...810..126C", "2015BAAA...57..203F", "2015MNRAS.451.3340K", "2016A&A...590A.110C", "2016ApJ...829...84C", "2016MNRAS.459.3460K", "2016MNRAS.460.2110F", "2016MNRAS.462.2033G", "2017A&A...601A..18B", "2017ASPC..510....3S", "2017ApJ...836..199H", "2017ApJ...846...16S", "2017ApJ...847..140C", "2017CRGeo.349..187A", "2017MNRAS.465..885S", "2017MNRAS.465.4611P", "2017MNRAS.470.2938F", "2018A&A...609A.123M", "2018A&A...611A..24H", "2018ApJ...854...88R", "2018ApJ...865...34C", "2018ApJ...869..115S", "2018MNRAS.474.5122K", "2018MNRAS.475.2632H", "2018MNRAS.480.3916M", "2018PASJ...70S..53I", "2019A&A...623A..16S", "2019A&A...632A..83S", "2019ApJ...872..207A", "2019ApJ...876..108D", "2019ApJS..240...26S", "2019FrASS...6....5H", "2019FrASS...6...51T", "2019MNRAS.485.3499C", "2019MNRAS.489.5326L", "2019MNRAS.490.2760K", "2019NewA...67...89N", "2019PASA...36...29A", "2020A&A...642A..76Z", "2020A&A...644A..27B", "2020ApJ...891...84C", "2020ApJ...891..168W", "2020ApJ...896...11F", "2020ApJ...898...67G", "2020ApJ...905..158W", "2020MNRAS.494.1971C", "2020MNRAS.494.3675C", "2020SSRv..216...70L", "2021A&A...649L..13D", "2021A&A...655A..76S", "2021ApJ...907...33Y", "2021ApJ...908...92H", "2021ApJ...914...87G", "2021ApJ...916...83A", "2021ApJ...918...39L", "2021ApJ...921...23G", "2021MNRAS.502.6188M", "2021MNRAS.503..768Y", "2021MNRAS.503.5006S", "2021PASJ...73S.385S", "2021RAA....21..188Z", "2021arXiv210200163M", "2022A&A...665A..77S", "2022AJ....164..257H", "2022ApJ...934..174I", "2022ApJ...935...57C", "2022ApJ...935...95G", "2022ApJ...940..106A", "2022ApJ...941..122C", "2022FrASS...9.0900B", "2022MNRAS.510.2644S", "2022MNRAS.513.1275A", "2022MNRAS.514.1575C", "2022MNRAS.517..885O", "2022MNRAS.517.1138S", "2022MNRAS.517.4940W", "2022MNRAS.517.5003B", "2023A&A...670A..38G", "2023A&A...673A..76S", "2023ASPC..534..193P", "2023ASPC..534..233P", "2023ApJ...946..114M", "2023ApJ...949...19M", "2023ApJ...955..113X", "2023MNRAS.519..285S", "2023MNRAS.521.4579S", "2023MNRAS.522.4972S", "2023MNRAS.525..364P", "2023arXiv231212504C", "2024ApJ...961..100A", "2024MNRAS.52710279C", "2024OJAp....7E..18H", "2024arXiv240408046H" ]
[ "astronomy" ]
9
[ "diffusion", "ISM: magnetic fields", "magnetohydrodynamics: MHD", "stars: formation", "turbulence", "Astrophysics - Solar and Stellar Astrophysics" ]
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[ "10.1088/0004-637X/785/1/69", "10.48550/arXiv.1403.0582" ]
1403
1403.0582_arXiv.txt
\label{sec:intro} The formation of stars begins with dense molecular cores \citep{2007ARA&A..45..565M,2009sfa..book..254A}. These cores form through the concentration of overdense regions within turbulent, filamentary GMCs; subsequent core collapse leads to protostellar (or protobinary)/disk systems. Magnetic fields are important at all scales during this process \citep{2007ARA&A..45..565M,2012ARA&A..50...29C}: the cloud-scale magnetic field can limit compression in interstellar shocks that create dense clumps and filaments in which cores form, while the local magnetic field within individual cores can prevent collapse if it is large enough \citep{1956MNRAS.116..503M,1966MNRAS.132..359S,1976ApJ...210..326M}, and can help to remove angular momentum during the disk formation process if cores are successful in collapsing \citep{1985prpl.conf..320M,1991ApJ...373..169M, 2003ApJ...599..363A, 2013PPVI...Li}. The significance of magnetic fields in self-gravitating cores can be quantified by the ratio of mass to magnetic flux; only if the mass-to-flux ratio exceeds a critical value is gravitational collapse possible. How the mass-to-flux ratio increases from the strongly-magnetized interstellar medium to weakly-magnetized stars is a fundamental problem of star formation \citep{1987ARA&A..25...23S, 2007ARA&A..45..565M}. Here, as suggested in \citet[][hereafter CO12]{2012ApJ...744..124C}, we consider core formation in GMCs with highly supersonic turbulence and non-ideal MHD. Magnetic fields are coupled only to charged particles, while the gas in GMCs and their substructures is mostly neutral. The ability of magnetic fields to affect core and star formation thus depends on the collisional coupling between neutrals and ions. Ambipolar diffusion is the non-ideal MHD process that allows charged particles to drift relative to the neutrals, with a drag force proportional to the collision rate \citep{1956pfig.book.....S}. Ambipolar drift modifies the dynamical effect of magnetic fields on the gas, and may play a key role in the star formation. In classical theory, quasi-static ambipolar diffusion is the main mechanism for prestellar cores to lose magnetic support and reach supercritical mass-to-flux ratios. Through ambipolar drift, the mass within dense cores can be redistributed, with the neutrals diffusing inward while the magnetic field threading the outer region is left behind \citep{1979ApJ...228..475M}. However, the quasi-static evolution model \citep[e.g.][]{1999osps.conf..305M,2001ApJ...547..272C} gives a prestellar core lifetime considerably longer (up to a factor of 10) than the gravitational free-fall timescale, $t_\mathrm{ff}$, while several observational studies have shown that cores only live for $(2-5)$ $t_\mathrm{ff}$ \citep[e.g.][]{2007prpl.conf...33W, 2009ApJS..181..321E}. The failure of the traditional picture to predict core lifetimes indicates that supercritical cores may not have formed quasi-statically through ambipolar diffusion. Indeed, it is now generally recognized that, due to pervasive supersonic flows in GMCs, core formation is not likely to be quasi-static. Realistic star formation models should take both ambipolar diffusion and large-scale supersonic turbulence into consideration. This turbulence may accelerate the ambipolar diffusion process \citep{2004ApJ...603..165H,2004ApJ...609L..83L}, with an analytic estimate of the enhanced diffusion rate by a factor of 2$-$3 for typical conditions in GMCs \citep{2002ApJ...570..210F}. In our previous work (\hyperlink{CO12}{CO12}), we investigated the physical mechanism driving enhanced ambipolar diffusion in one-dimensional C-type shocks. These shocks pervade GMCs, and are responsible for the initial compression of gas above ambient densities. We obtained a formula for the C-shock thickness as a function of density, magnetic field, shock velocity, and ionization fraction, and explored the dependence of shock-enhanced ambipolar diffusion on environment through a parameter study. Most importantly, we identified and characterized a transient stage of rapid ambipolar diffusion at the onset of shock compression, for one-dimensional converging flows. For an interval comparable to the neutral-ion collision time and before the neutral-ion drift reaches equilibrium, the neutrals do not experience drag forces from the ions. As a consequence, the initial shock in the neutrals is essentially unmagnetized, and the neutrals can be very strongly compressed. This transient stage, with timescale $t_\mathrm{transient}\sim 1$~Myr (but depending on ionization), can create dense structures with much higher $\rho/B$ than upstream gas. \hyperlink{CO12}{CO12} suggested this could help enable supercritical core formation. \hyperlink{CO12}{CO12} also found that (1) the perpendicular component of the magnetic field is the main determinant of the shock compression, and (2) the perpendicular component of the magnetic field $B_\perp$ must be weak ($\lesssim 5~\mu$G) for transient ambipolar diffusion in shocks to significantly enhance $\rho/B_\perp$. Observations of nearby clouds provide direct constraints on the role of magnetic fields, as well as other properties of prestellar cores. The typical mean mass-to-flux ratio of dark cloud cores is $\Gamma\sim 2$ (in units of critical value; see Equation~(\ref{GammaDef})) from Zeeman studies \citep{2008A&A...487..247F,2008ApJ...680..457T}. Due to the instrumental limitations, magnetic field observations in solar-mass and smaller scale regions are relatively lacking compared with observations of larger scales \citep[see review in][]{2012ARA&A..50...29C}, however. Surveys in nearby clouds have found that prestellar cores have masses between $\sim 0.1-10$~M$_\odot$ and sizes $\sim 0.01-1$~pc \citep[e.g.][]{2001A&A...372L..41M, 2009ApJ...691.1560I, 2009ApJ...699..742R,2013MNRAS.432.1424K}. In addition, a mass-size relation has been proposed as a power law $M\propto R^k$, with $k=1.2-2.4$ dependent on various molecule tracers \citep[e.g.][]{1996ApJ...471..816E,2010MNRAS.402..603C,2010ApJ...723..492R,2013MNRAS.432.1424K}. The magnetic field strength within prestellar cores is important for late evolution during core collapse, since disk formation may be suppressed by magnetic braking (for recent simulations see \citealt{2003ApJ...599..363A,2008A&A...477....9H,2008ApJ...681.1356M,2011A&A...528A..72H}; or see review in \citealt{2013PPVI...Li}). However, many circumstellar disks and planetary systems have been detected \citep[e.g.][]{2001ApJ...553L.153H,2010A&A...512A..40M}, suggesting that the magnetic braking ``catastrophe" seen in many simulations does not occur in nature. The proposed solutions include the misalignment between the magnetic and rotation axes \citep[e.g.][]{2009A&A...506L..29H,2010MNRAS.409L..39C,2012A&A...543A.128J,2013ApJ...767L..11K}, turbulent reconnection and other turbulent processes during the rotating collapse \citep[e.g.][]{2012ApJ...747...21S,2012MNRAS.423L..40S,2013MNRAS.432.3320S}, and non-ideal MHD effects including ambipolar diffusion, Hall effect, and Ohmic dissipation \citep[e.g.][]{2010ApJ...716.1541K,2011ApJ...738..180L,2011PASJ...63..555M,2012A&A...541A..35D,2013ApJ...763....6T}. If prestellar cores have sufficiently weak magnetic fields, however, braking would not be a problem during disk formation \citep[e.g.][]{2008ApJ...681.1356M,2013PPVI...Li,2013ApJ...774...82L}. Therefore, the magnetic field (and mass-to-flux ratio) within a prestellar core is important not just for the ability of the core to collapse, but also of a disk to form. Fragmentation of sheetlike magnetized clouds induced by small-amplitude perturbation and regulated by ambipolar diffusion has been widely studied \citep[e.g.][]{2000ApJ...532..361I,2004ApJ...607L..39B,2005ApJ...622..393B,2006ApJ...652..442C,2009NewA...14..221B}. Analogous fully three-dimensional simulations have also been conducted \citep[e.g.][]{2007MNRAS.380..499K}. Supercritical cores formed in the flattened layer have masses $\sim 0.1-10$~M$_\odot$ \citep[e.g.][]{2000ApJ...532..361I,2009NewA...14..221B}, at timescales $\sim 1-10$~Myr dependent on the initial mass-to-flux ratio of the cloud \citep[e.g.][]{2000ApJ...532..361I,2007MNRAS.380..499K,2009NewA...14..221B}. The above cited simulations start from relatively high densities ($\sim 10^4$~cm$^{-3}$; e.g. \citealt{2007MNRAS.380..499K}) and included only the low-amplitude perturbations. Alternatively, \citet{2004ApJ...609L..83L} and \citet{2005ApJ...631..411N} took the formation of these overdense regions into consideration by including a direct treatment of the large-scale supersonic turbulence. They demonstrated that ambipolar diffusion can be sped up locally by the supersonic turbulence, forming cores with masses $\sim 0.5$~M$_\odot$ and sizes $\sim 0.1$~pc within $\sim 2$~Myr, while the strong magnetic field keeps the star formation efficiency low ($1-10\%$). Similarly, \citet{2009NewA...14..483B} found that turbulence-accelerated, magnetically-regulated core formation timescales are $\sim 1$~Myr in two-dimensional simulations of magnetized sheet-like clouds, with corresponding three-dimensional simulations showing comparable results \citep{2008ApJ...679L..97K,2011ApJ...728..123K}. In addition, \citet{2008ApJ...687..354N} measured the core properties in their three-dimensional simulations to find $L_\mathrm{core}\sim 0.04-0.14$~pc, $\Gamma_\mathrm{core}\sim 0.3-1.5$, and $M_\mathrm{core}\sim 0.15-12.5$~M$_\odot$, while \citet{2009NewA...14..483B} found a broader core mass distribution $M_\mathrm{core}\sim 0.04-25$~M$_\odot$ in their parameter study using thin-sheet approximation. Supersonic turbulence within GMCs extends over a wide range of spatial scales \citep{2004RvMP...76..125M,2007prpl.conf...63B}. Although turbulence contains sheared, diverging, and converging regions in all combinations, regions in which there is a large-scale convergence in the velocity field will strongly compress gas, creating favorable conditions for the birth of prestellar cores. \citet[][hereafter GO11]{2011ApJ...729..120G} investigated core formation in an idealized model containing both a large-scale converging flow and multi-scale turbulence. These simulations showed that the time until the first core collapses depends on inflow Mach number ${\cal M}$ as $t_\mathrm{collapse}\propto {\cal M}^{-1/2}$. With a parameter range ${\cal M} = 1.1$ to $9$, cores formed in the \hyperlink{GO11}{GO11} simulations had masses $0.05-50$~M$_\odot$. Following similar velocity power spectrum but including ideal MHD effects, \citet{2014arXiv1401.6096M} performed simulations with sink particle, radiative transfer, and protostellar outflows to follow the protostar formation in turbulent massive clump. They demonstrated that the median stellar mass in the simulated star cluster can be doubled by the magnetic field, from $0.05$~M$_\odot$ (unmagnetized case) to $0.12$~M$_\odot$ (star cluster with initial mass-to-flux ratio $\Gamma=2$). This is qualitatively consistent with the conclusion in \citet{2013ApJ...774L..31I}, that the mass of the cores formed in the post-shock regions created by cloud-cloud collision is positively related to (and dominated by) the strong magnetic field in the shocked layer. Note that, though the main focus of \citet{2013ApJ...774L..31I} is the cloud's ability to form massive cores ($\sim 20-200$~M$_\odot$ in their simulations), the idea of cloud-cloud collision is very similar to the converging flows setup adopted in \hyperlink{GO11}{GO11} and this study. In this paper, we combine the methods of \hyperlink{CO12}{CO12} for modeling ambipolar diffusion with the methods of \hyperlink{GO11}{GO11} for studying self-gravitating structure formation in turbulent converging flows. Our numerical parameter study focuses on the level of ambipolar diffusion (controlled by the ionization fraction of the cloud) and the obliquity of the shock (controlled by the angle between the magnetic field and the upstream flow). We show that filamentary structures similar to those seen in observations \citep[see review in][]{2013PPVI...Andre} develop within shocked gas layers, and that cores form within these filaments. We measure core properties to test their dependence on these parameters. As we shall show, our models demonstrate that low-mass supercritical cores can form for all magnetic obliquities and all levels of ionization, including ideal MHD. However, our models also show that ambipolar diffusion affects the magnetization of dynamically-formed cores. The outline of this paper is as follows. We provide a theoretical analysis of oblique MHD shocks in Section~\ref{theory}, pointing out that a quasi-hydrodynamic compression ratio (which is $\sim 5$ times stronger than in \textit{fast} MHD shocks for the parameters we study) can exist when the converging flow is nearly parallel to the magnetic field. We also show that shock compression cannot increase the mass-to-flux ratio except in the nearly-parallel case or with ambipolar diffusion. Section~\ref{sec: methods} describes methods used in our numerical simulations and data analysis, including our model parameter set and method for measuring magnetic flux within cores. The evolution of gas structure (including development of filaments) and magnetic fields for varying parameters is compared in Section~\ref{sec::evolution}. In Section~\ref{sec:results} we provide quantitative results for masses, sizes, magnetizations, and other physical properties of the bound cores identified from our simulations. Implications of these results for core formation is discussed in Section~\ref{sec::CF}, where we argue that the similarity of core masses and sizes among models with different magnetizations and ionizations can be explained by anisotropic condensation preferentially along the magnetic field. Section~\ref{sec: summary} summarizes our conclusions.
\label{sec: summary} In this work, we have used numerical simulations to study core formation in magnetized, highly dynamic environments, including the effect of ambipolar diffusion. Our simulations are fully three-dimensional, including a large-scale convergent flow, local turbulence, and self-gravity, and allow for varying ambipolar diffusion levels (parameterized by the ionization fraction coefficient $\chi_{i0}$) and shock obliquity (parameterized by the angle $\theta$ between the converging inflow and the global magnetic field). Filaments and then cores form in post-shock dense layers, with dense structures very similar to those found in observations. In all of our models (with or without ambipolar diffusion), magnetically supercritical cores form with physical properties similar to those found in observations. However, our parameter survey suggests that the transient ambipolar diffusion timescale and \textit{quasi-hydrodynamic} shocks are crucial in setting the magnetization of cores formed in post-shock regions. In addition, we demonstrate and quantitatively explain how low-mass supercritical cores form in strongly-magnetized regions, via anisotropic condensation along the magnetic field. Our main conclusions are as follows: \begin{enumerate} \item Under typical GMC conditions, isotropic formation of low-mass supercritical cores is forbidden under ideal MHD by the relatively strong magnetic support (Equation~(\ref{McritSPH})). This is true even downstream from strong MHD shocks where gas density is enhanced, because the magnetic field is compressed as well. In fact, for a spherical volume of given mass, the mass-to-flux ratio is generally larger for pre-shock conditions than post-shock conditions (Equation~(\ref{GammaComp}); except for the special case described in \#2 below). For typical conditions, the minimum post-shock critical mass for a spherical volume exceeds $10$~M$_\odot$ when ideal MHD applies (Table~\ref{modelpar}, \ref{simresults}). This suggests that either transient ambipolar diffusion in shocks must be taken into consideration, or that core formation is not spherically symmetric. \item When the incoming flows are almost parallel to the background magnetic field, MHD shocks will have compound post-shock conditions, including the regular \textit{fast} mode \citep{1992phas.book.....S} and the \textit{quasi-hydrodynamic} mode in which gas is compressed more strongly (Figure~\ref{rfcomp}). This happens when the angle $\theta$ between the inflow and the magnetic field is smaller than a critical value, $\theta_\mathrm{crit}$ (Equation~(\ref{thetaCrit})). For small $\theta$, the post-shock layer will have relatively high gas density and weak magnetic field compared to \textit{fast}-mode MHD shocks (Table~\ref{simresults}). \item Our three-dimensional simulations demonstrate the effect of transient ambipolar diffusion, as earlier identified and explained by \hyperlink{CO12}{CO12}. During the earliest stage of shock formation ($t\lesssim 0.3$~Myr), a thin but extremely dense layer appears in the middle of the shocked region in models with ambipolar diffusion (Figure~\ref{Evo1} and \ref{Evo2}), just like the central dense peak in the one-dimensional shocks analyzed by \hyperlink{CO12}{CO12}. Consequently, post-shock densities are generally higher in models with lower ionizations (smaller $\chi_{i0}$; see Table~\ref{simresults}), which correspond to stronger ambipolar diffusion as predicted in \hyperlink{CO12}{CO12}. \item The ionization fraction is the main parameter controlling the transient ambipolar diffusion timescale needed for the gas to reach steady post-shock conditions ($t_\mathrm{transient}$). Models with smaller $\chi_{i0}$ have longer transient timescales (Equation~(\ref{t_tran})), indicating lower growth rate of the post-shock magnetic field and more weakly magnetized post-shock layers (Table~\ref{simresults}). Therefore, transient ambipolar diffusion is crucial in reducing the magnetic support in the post-shock regions (see $M_\mathrm{mag,sph}$ and $R_\mathrm{mag,sph}$ in Table~\ref{simresults}). \item The filament network in more strongly magnetized post-shock cases is similar to those found in observations: in addition to large-scale main filaments, there are many thinner, less-prominent sub-filaments parallel to the magnetic field \citep{2008ApJ...680..428G,2011ApJ...734...63S,2012A&A...543L...3H,2013A&A...550A..38P,2013PPVI...Andre}. Dense cores form within the large-scale main filaments for all models. \item In our simulations, magnetically supercritical cores are able to form in the shock-compressed dense layers in all models, and the first collapse occurs at $t \lesssim 0.6$~Myr in most cases. Cores formed in our simulations have masses $\sim 0.04-2.5$~M$_\odot$ and sizes $\sim 0.015-0.07$~pc (Table~\ref{modelsum} and Figure~\ref{Dist}), similar to the values obtained in observations \citep[e.g.][]{2001A&A...372L..41M, 2009ApJ...691.1560I, 2009ApJ...699..742R,2013MNRAS.432.1424K}. The medians from the distributions are $0.47$~M$_\odot$ and $0.03$~pc. The mass-size relationship derived from our cores, $M\propto L^{2.3}$, also agrees with observations \citep[e.g.][]{1996ApJ...471..816E,2010MNRAS.402..603C,2010ApJ...723..492R,2013MNRAS.432.1424K}. \item Our results show that the core mass and size are relatively independent of both the ambipolar diffusion and the upstream magnetic obliquity (Figure~\ref{MaModel}). Hydrodynamic and ideal MHD models also have very similar core masses and sizes. The core masses for ideal MHD cases with oblique shocks are more than an order of magnitude lower than the magnetic critical mass for a spherical region in the post-shock environment. Thus, simple estimates of the form in Equation~(\ref{McritSPH}) should not be used in predicting magnetically supercritical core masses from ambient environmental conditions in a GMC. \item The magnetic field of cores follows the same trends as the post-shock magnetization, in terms of variation with the upstream magnetic obliquity and ionization (Table~\ref{simresults}, \ref{modelsum}). This indicates that further ambipolar diffusion is limited during the core building phase, and instead cores form by anisotropic self-gravitating contraction as described in Section~\ref{sec::CF}. The mass-to-flux ratio in cores secularly increases with decreasing ionization (Figure~\ref{MaModel}), ranging from $\Gamma \sim 0.5$ to $7.5$ (Figure~\ref{GaDist}). From all models combined, the median mass-to-flux ratio within cores is $\Gamma\sim 3$ (Figure~\ref{GaDist}), agreeing with the observed range of $\Gamma$ \citep[$\Gamma\sim 1-4$;][]{2008A&A...487..247F,2008ApJ...680..457T}. \item Anisotropic self-gravitating condensation is likely the dominant mechanism for supercritical core formation in magnetized environments, regardless the magnetization strength and ionization fraction. Figures~\ref{coreEvo} and \ref{Vtime} clearly show how gas preferentially flows along the magnetic field lines in all models, creating dense cores that are both magnetically and thermally supercritical. The theoretical analysis of Section~\ref{anisotropic} shows that the characteristic mass expected from anisotropic contraction (Equation~(\ref{McritNew})) is similar to the median core mass obtained from our simulations (Figure~\ref{Dist}). For anisotropic core formation in a post-shock region, the critical mass is expected to depend only on the momentum flux entering the shock. We believe this explains why core masses in our simulations are similar regardless of the ionization level, whether the converging flow is nearly parallel to or highly oblique to the upstream magnetic field, or indeed whether the medium is even magnetized at all. \end{enumerate}
14
3
1403.0582
We investigate the roles of magnetic fields and ambipolar diffusion during prestellar core formation in turbulent giant molecular clouds, using three-dimensional numerical simulations. Our simulations focus on the shocked layer produced by a converging large-scale flow and survey varying ionization and the angle between the upstream flow and magnetic field. We also include ideal magnetohydrodynamic (MHD) and hydrodynamic models. From our simulations, we identify hundreds of self-gravitating cores that form within 1 Myr, with masses M ~ 0.04-2.5 M <SUB>⊙</SUB> and sizes L ~ 0.015-0.07 pc, consistent with observations of the peak of the core mass function. Median values are M = 0.47 M <SUB>⊙</SUB> and L = 0.03 pc. Core masses and sizes do not depend on either the ionization or upstream magnetic field direction. In contrast, the mass-to-flux ratio does increase with lower ionization, from twice to four times the critical value. The higher mass-to-flux ratio for low ionization is the result of enhanced transient ambipolar diffusion when the shocked layer first forms. However, ambipolar diffusion is not necessary to form low-mass supercritical cores. For ideal MHD, we find similar masses to other cases. These masses are one to two orders of magnitude lower than the value M <SUB>mag, sph</SUB> = 0.007B <SUP>3</SUP>/(G <SUP>3/2</SUP>ρ<SUP>2</SUP>) that defines a magnetically supercritical sphere under post-shock ambient conditions. This discrepancy is the result of anisotropic contraction along field lines, which is clearly evident in both ideal MHD and diffusive simulations. We interpret our numerical findings using a simple scaling argument that suggests that gravitationally critical core masses will depend on the sound speed and mean turbulent pressure in a cloud, regardless of magnetic effects.
false
[ "upstream magnetic field direction", "magnetic field", "magnetic fields", "post-shock ambient conditions", "low ionization", "lower ionization", "turbulent giant molecular clouds", "prestellar core formation", "Core masses", "magnetic effects", "varying ionization", "field lines", "enhanced transient ambipolar diffusion", "similar masses", "ambipolar diffusion", "masses", "diffusive simulations", "turbulent pressure", "low-mass supercritical cores", "ideal MHD" ]
11.187957
10.586747
183
746034
[ "Guzman-Ramirez, L.", "Lagadec, E.", "Jones, D.", "Zijlstra, A. A.", "Gesicki, K." ]
2014MNRAS.441..364G
[ "PAH formation in O-rich planetary nebulae" ]
33
[ "European Southern Observatory, Alonso de Córdova 3107, Casilla 19001, Santiago, Chile; Jodrell Bank Centre for Astrophysics, School of Physics and Astronomy, University of Manchester, Manchester M13 9PL, UK", "Astronomy department, Cornell University, Ithaca, NY 14853-6801, USA", "European Southern Observatory, Alonso de Córdova 3107, Casilla 19001, Santiago, Chile; Universidad de Atacama, Departamento de Física, Copayapu 485, Copiapó, Chile", "Jodrell Bank Centre for Astrophysics, School of Physics and Astronomy, University of Manchester, Manchester M13 9PL, UK", "Centre for Astronomy, Faculty of Physics, Astronomy and Informatics, Nicolaus Copernicus University, Grudziadzka 5, PL-87-100 Torun, Poland" ]
[ "2014A&A...568A..95O", "2015ASPC..497..155G", "2015ApJ...806...86D", "2015EAS....71..259G", "2015MNRAS.449L..56D", "2015MNRAS.451L...1G", "2015MNRAS.452.2911A", "2015RMxAA..51..221Z", "2016A&A...593A..92B", "2016AJ....151...93O", "2016JPhCS.728c2011C", "2016MNRAS.456L..89C", "2017ApJS..231...22O", "2017IAUS..323..121S", "2017IAUS..323..141Z", "2017MNRAS.470..626M", "2017MNRAS.470.3250J", "2017MNRAS.471.3529S", "2017PASP..129h2001P", "2018MNRAS.473.4476G", "2018MNRAS.475.2282V", "2018MNRAS.480.2423I", "2018arXiv180608244J", "2019MNRAS.485.3360T", "2019MNRAS.488..451M", "2019ibfe.book.....B", "2020Galax...8...44L", "2021AJ....162..248H", "2021MNRAS.508.2254R", "2022ApJ...927...46S", "2022ApJ...933L..24R", "2022arXiv220700547Z", "2023MNRAS.520.1073I" ]
[ "astronomy" ]
5
[ "circumstellar matter", "planetary nebulae: general", "infrared: stars", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1984A&A...137L...5L", "1988A&A...202L..27A", "1991A&A...243L...9Z", "1993ApJ...413..641V", "1998A&A...331L..61W", "1998Natur.391..868W", "1999A&A...352..587S", "1999ApJ...513L.135C", "2000A&A...358.1058G", "2002A&A...384..620A", "2002MNRAS.332..879C", "2003A&A...403..659A", "2003MNRAS.344..501P", "2003MNRAS.345.1291R", "2004Msngr.117...12L", "2005A&A...429..977W", "2005ARA&A..43..435H", "2006A&A...451..925G", "2007A&A...463..251U", "2008ApJ...680.1206G", "2008MNRAS.391L.113G", "2009A&A...495L...5P", "2009A&A...505..249M", "2010A&A...521L..60B", "2010MNRAS.408.2312J", "2010MNRAS.409..881P", "2011MNRAS.410.1349C", "2011MNRAS.414.1667G", "2011MNRAS.417...32L", "2011RMxAA..47...83P", "2012A&A...541A.112K", "2012EPJWC..1906009U", "2012MNRAS.422.1804T", "2013MNRAS.435..975R" ]
[ "10.1093/mnras/stu454", "10.48550/arXiv.1403.1856" ]
1403
1403.1856_arXiv.txt
All low and intermediate-mass stars, with initial masses between 0.8 and 8 M$_{\odot}$, end their life as a white dwarf (WD). Some of them will experience the planetary nebulae (PNe) phase before the end of their lives, when the ultraviolet emission from the hot WD photodissociates and photoionises the expanding gas and dust that was ejected in the previous phase, when the star is on the asymptotic giant branch (AGB). During the AGB phase the intense mass loss (from $10^{-6}$ to $10^{-4}\,\rm M_{\odot}\, yr^{-1}$) leads to the formation of a circumstellar envelope made of gas and dust \citep{kwok}. The dust is seen in emission as an infrared excess superposed on the stellar spectral energy distribution (SED). In the AGB phase, a star may evolve from being oxygen-rich to being carbon-rich. The change occurs when carbon produced by He-burning is brought to the surface by dredge-up processes in the stellar interior thereby increasing the C/O ratio until it exceeds unity and a carbon star is formed \citep{herwig}. This process depends on stellar mass: \cite{vassi} showed the third dredge-up occurs if the core mass of the star $M_{cs} > 1.5$\,M$_{\odot}$. This predicts a clear distinction, with some (lower mass) stars showing oxygen-rich and some (higher mass) stars carbon-rich ejecta. In the molecular ejecta, the CO molecule locks away the less abundant element, leaving the remaining free O or C to drive the chemistry and dust formation. Oxygen-rich shells are characterised by silicate dust. Amorphous silicates dominate the 7-25$\upmu$m region in the IR spectrum, the 9.7$\upmu$m feature and the 18$\upmu$m are very strong bands that occur in emission or (self-)absorption. Crystalline silicate emission features occur near 28, 33 and 43$\upmu$m \citep{sylvester}. Carbon-rich shells show Polycyclic Aromatic Hydrocarbon (PAH) emission bands and carbonaceous amorphous dust. The PAH emission bands are at 6.2, 7.7, 8.6, and 11.3$\upmu$m \citep{leger84}. \cite{zijlstra91} observed the first evidence for mixed chemistry when one planetary nebula (PN) (IRAS07027-7934) with strong PAH emission bands, was found to also show a 1.6GHz OH maser line. Observations made with the Infrared Space Observatory ({\it ISO}) uncovered several further cases, where PAH emission in PNe occurred together with emission bands of silicates usually found in O-rich shells (Waters et al. 1998a,b; Cohen et al., 1999,2002)\nocite{cohen99,cohen02}. These PNe were shown to be the envelope of late/cool [WC] type stars and the mixed chemistry was therefore explained as an evolutionary change in the central star - a recent thermal pulse, which made the star turn from O-rich to C-rich. This is the so-called very late thermal pulse (VLTP). In all these objects the torus remained O-rich, while the C-rich material is observed in the outflows. \cite{guten} and \cite{perea09} showed that the mixed chemistry phenomenon is widespread amongst Galactic bulge (GB) PNe. Their {\it Spitzer} observations show that the simultaneous presence of O and C-rich dust features is common, and is not restricted to objects with late/cool [WC] type stars. The traditional explanation relating mixed chemistry to a recent evolution towards carbon-star is highly unlikely for Bulge objects, as these old, low-mass stars are not expected to show substantial third dredge-up, and therefore should not show enhanced C/O ratios. The few AGB carbon stars in the Bulge do not originate from third dredge-up \citep{azzo}. \cite{stephan} found that 4 AGB stars out of a sample of 27 objects in the GB show Tc, evidence of a third dredge-up. However, they also observed 45 AGB stars in the GB finding they are all O-rich \citep{stephan2}. The third dredge-up is thus occurring in some of these objects, but is not large enough to form carbon-rich AGB stars. \cite{me11} analysed a sample of 40 GBPNe, with the mixed chemistry phenomenon found in 30 nebulae. {\it HST} images and UVES spectra showed that the mixed chemistry is not related to the presence of emission-line stars, as it is in the Galactic disk population. Instead, a strong correlation is found with morphology, and the presence of a dense dusty equatorial structure (torus). The mixed chemistry phenomenon occurring in the GBPNe is best explained through hydrocarbon chemistry in an UV-irradiated, dense torus. One way to test this theory is to spatially resolve the PAHs in these PNe. If the PAHs are only present in the outflows, this would support the hypothesis that they originate from a VLTP, implying that the central star changed from O-rich to C-rich. On the other hand, if the PAHs are concentrated in the torus, this would be more consistent with their formation resulting from the photodissociation of CO, meaning that the central star does not have to experience a third dredge-up nor VLTP to be able to produce C-rich molecules. Aiming to detect the PAHs either in the torus or the outflows, we selected 11 targets to be observed using the {\it VLT} spectrometer and imager for the mid--infrared (VISIR) instrument on the Very Large Telescope ({\it VLT}). We obtained images in three filters, PAH1 (8.59$\upmu$m), PAH2 (11.25$\upmu$m) and SIV (10.49$\upmu$m). We also used the long-slit spectrograph to analyse in more detail a sub-sample of three objects. This paper is organised as follows. In Section 2, we present the observing technique. In Section 3, we show the results of the observations of 11 PNe, presenting the acquired images and spectra. In Section 4, we discuss the main results before presenting the final conclusions in Section 5.
The main results and conclusions of this paper can be summarised as follows: \begin{itemize} \item Emission of the PAH features at 8.6$\upmu$m and 11.3$\upmu$m and the [SIV] line were imaged using the VISIR instrument in the {\it VLT}. We find dense, dusty, toroidal structures in 10 (out of 11) objects we observed. We detected PAHs in the torus of 8 of them (we could not resolved the torus in Th3-4 and H1-43), this finding confirms the proposed PAH formation scenario whereby CO is photodissociated as described by the models presented in \citet{me11}. \item For most of the objects the [SIV] line shows emission in an inner region than the PAHs. One would expect these emission lines to come from an inner region, where the elements can recombine, while PAHs would form just outside this region, where some UV-photons penetrate and dissociate CO. The free C then will aggregate in these regions, leading to the formation of the PAHs. \item We found that the 12.4$\upmu$m H$_2$ line anti-correlates with the temperature of the central star (i.e. PNe with hotter central stars, display lower levels of H$_2$ emission). This shows that the H$_2$ is being photodissociated by the UV photons of the central star. \end{itemize}
14
3
1403.1856
Polycyclic aromatic hydrocarbons (PAHs) have been observed in O-rich planetary nebulae towards the Galactic bulge. This combination of oxygen-rich and carbon-rich material, known as dual-dust or mixed chemistry, is not expected to be seen around such objects. We recently proposed that PAHs could be formed from the photodissociation of CO in dense tori. In this work, using VISIR/VLT, we spatially resolved the emission of the PAH bands and ionized emission from the [S IV] line, confirming the presence of dense central tori in all the observed O-rich objects. Furthermore, we show that for most of the objects, PAHs are located at the outer edge of these dense/compact tori, while the ionized material is mostly present in the inner parts of these tori, consistent with our hypothesis for the formation of PAHs in these systems. The presence of a dense torus has been strongly associated with the action of a central binary star and, as such, the rich chemistry seen in these regions may also be related to the formation of exoplanets in post-common-envelope binary systems.
false
[ "dense central tori", "dense tori", "such objects", "ionized emission", "post-common-envelope binary systems", "PAHs", "the rich chemistry", "a central binary star", "exoplanets", "Galactic", "these dense/compact tori", "oxygen-rich and carbon-rich material", "the ionized material", "a dense torus", "CO", "all the observed O-rich objects", "dual-dust or mixed chemistry", "these tori", "the Galactic bulge", "PAH" ]
10.572404
10.665863
-1
483053
[ "Pyo, Tae-Soo", "Hayashi, Masahiko", "Beck, Tracy L.", "Davis, Christopher J.", "Takami, Michihiro" ]
2014ApJ...786...63P
[ "[Fe II] Emissions Associated with the Young Interacting Binary UY Aurigae" ]
8
[ "Subaru Telescope, National Astronomical Observatory of Japan, 650 North A'ohoku Place, Hilo, HI 96720, USA", "National Astronomical Observatory of Japan, 2-21-1 Osawa, Mitaka, Tokyo 181-8588, Japan; School of Mathematical and Physical Science, The Graduate University for Advanced Studies (SOKENDAI), Hayama, Kanagawa 240-0193, Japan", "Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21218, USA", "Astrophysics Research Institute, Liverpool John Moores University, Liverpool Science Park, 146 Brownlow Hill, Liverpool L3 5RF, UK", "Institute of Astronomy and Astrophysics, Academia Sinica, P.O. Box 23-141, Taipei 10617, Taiwan" ]
[ "2014ApJ...792...56S", "2014ApJ...793...10T", "2015PKAS...30..109P", "2016A&ARv..24....5D", "2019ApJ...884..159B", "2020AJ....160...39U", "2020ApJ...897L..22S", "2024ApJ...963..159P" ]
[ "astronomy" ]
4
[ "ISM: jets and outflows", "stars: formation", "stars: individual: UY Aurigae", "stars: low-mass", "stars: pre-main sequence", "techniques: high angular resolution", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1944PASP...56..123J", "1983ApJ...268..368E", "1991A&A...248..485D", "1991ApJ...370L..35A", "1992ApJ...396..178F", "1993ApJ...408L..49R", "1994A&A...286..149D", "1994ARA&A..32..465M", "1995ApJ...444L..93H", "1996ApJ...467L..77A", "1997A&AS..126..437H", "1997ASPC..121..792R", "1997ApJ...480..741K", "1998A&A...331..977K", "1998A&A...332..867D", "1998ApJ...499..883C", "1998ApJ...499L..75F", "1998ApJ...508..186K", "1998Natur.395..355R", "1999A&A...351..954D", "2000AJ....120.3177R", "2000ApJ...535..325W", "2000PASJ...52...81I", "2000prpl.conf..759K", "2000prpl.conf..789S", "2001IAUS..200.....Z", "2001MNRAS.326..524D", "2002A&A...387..550G", "2002A&A...393.1035N", "2002ApJ...568L..53T", "2002ApJ...570..724P", "2003A&A...397..675T", "2003A&A...397..693D", "2003AJ....126..327R", "2003AJ....126.2009B", "2003ApJ...583..334H", "2003ApJ...590..340P", "2003ApJ...599L..41E", "2003SPIE.4841.1581M", "2004A&A...415..189M", "2004A&A...417..247W", "2004AJ....127.1747H", "2005ApJ...618..817P", "2005ApJ...625L.131D", "2005ApJ...632L.135M", "2005prpl.conf.8174M", "2006ApJ...641..357T", "2006ApJ...646..319E", "2006ApJ...649..836P", "2007AJ....134..880H", "2007ApJ...657..897K", "2007ApJ...670L..33T", "2007prpl.conf..379D", "2007prpl.conf..395M", "2008A&A...478..453M", "2008A&A...487.1019G", "2008AJ....135.2496C", "2008AJ....135.2526C", "2008ApJ...676..472B", "2008ApJ...678.1109M", "2008ApJ...681..391M", "2008ApJ...687.1117F", "2009ApJ...694..582H", "2009ApJ...694..654P", "2009ApJ...704L..10M", "2009ApJS..185..289R", "2010A&A...511A...5G", "2010ApJ...708..485H", "2010ApJ...708L...5M", "2010ApJ...711.1280S", "2010ApJ...722.1360B", "2010AstL...36..653B", "2010Sci...327..306M", "2011A&A...528A...3D", "2011Ap&SS.335..125F", "2011MNRAS.411.2383K", "2012MNRAS.421.2025K", "2012MNRAS.427.2636K", "2013ApJ...768..110C" ]
[ "10.1088/0004-637X/786/1/63", "10.48550/arXiv.1403.3474" ]
1403
1403.3474_arXiv.txt
\label{sec:intro} UY Aur is a close binary system composed of classical T Tauri stars separated by 0$\farcs$89 \citep{Close1998,Duchene1999, Hioki2007}. \citet{Joy1944} was the first to identify the binarity of UY Aur. The secondary source UY Aur B is an infrared companion with a large extinction; \textit{A}$_V =~$22 - 12 magnitude \citep{Koresko1997}. The spectral types of the primary and secondary sources are M0 and M2.5, respectively \citep{HK2003}. In the optical forbidden emission lines of \OI\ and \SII , \citet{Hirth1997} identified a redshifted jet (HH 386) extending over a few arcseconds along a position angle (PA) of $\sim$~220$\degree$. They also reported that a blueshifted jet was evident in 1992 December. The driving source of the jets was not clear in their data, however; UY Aur A, B, or both sources. In their wide-field survey \citet{McGroarty2004} have since failed to detect the optical jets beyond what was discovered by Hirth et al. Molecular hydrogen emission has also been detected from UY Aur, though only from the secondary source, the infrared companion UY Aur B. This emission was interpreted as arising from accretion shocks associated with the circumstellar disk \citep{Herbst1995}. The circumbinary disk around UY Aur is the second such disk that has been resolved and imaged after the GG Tau A system. It was resolved for the first time in a millimeter interferometer survey \citep{Dutrey1994}. \citet{Duvert1998} found that the $^{13}$CO gas in the circumbinary disk is in Keplerian rotation and estimated the total mass of the binary as $\sim$~1.2 M$_\sun$. Using orbital parameters derived from observations obtained as far back as 1944, \citet{Close1998} estimated the total mass of the binary to be 1.61$^{+0.47}_{-0.67}$ $M_\sun$. The PA of the semi-major axis of the circumbinary disk is 135$\degree$; the inclination angle of the disk is $\sim$~42$\degree \pm$~3$\degree$ with respect to the line of sight. High spatial resolution near infrared imaging with adaptive optics (AO) has also revealed inner cavity betwen the circumstellar and circumbinary disks as well as clumpy structure in the circumbinary disk \citep{Close1998, Hioki2007}. The inner cavity or gap (i.e. the region with relatively low density) is produced by a continuous transfer of angular momentum from the binary to the outer circumbinary disk \citep{Artymowicz1991}. Indeed, \citet{Artymowicz1996} and \citet{Rozyczka1997} showed that material in the outer disk can penetrate the inner gap and episodically accrete onto the lower mass secondary star, provided the material has sufficiently high viscosity and temperature. Simulations of the UY Aur system by \citet{Gunther2002} suggest that the mass accretion rate from the outer circumbinary disk onto the secondary star might be five times higher and more phase dependent than that to the primary star. However, \citet{Hanawa2010} conversely showed that usually the mass accretion rate of the primary is higher than that of the secondary, and in late times the two accretion rates become similar. Many simulations produce a gas stream bridge between the circumstellar disks, as well as an accretion flow toward the individual circumstellar disks from the outer circumbinary disk \citep[e.g.][]{Fateeva2011}. \citet{Hanawa2010} showed that the bridge contains a strong shock front caused by the collision of opposing flows. Such a bridge structure has been detected in high-resolution coronagraphic H-band imaging of SR24 \citep{Mayama2010}. Many young stars are born in binary or multiple systems \citep[see e.g. the reviews by][]{Mathieu1994,Zinnecker2001,Duchene2007}. The binary frequency for solar type main-sequence stars is about 40 \%$-$60 \% \citep{Duquennoy1991, Fischer1992}, similar to the frequency (48.9 \%~$\pm$~5.3 \%) observed for the young stars in the Taurus molecular cloud \citep{Kohler1998}. Other nearby star forming regions show a binary freqency of 9 \%-32 \% \citep[][and references therein]{King2012a, King2012b}. Recent studies suggest that binarity is common in embedded protostars \citep{Haisch2004,Connelley2008a,Connelley2008b,Chen2013}. Jets or outflows have been observed from a few tens of multiple low-mass young stars \citep{Reipurth1993,Reipurth2000,Takami2003,Murphy2008,Mundt2010}. For single stars, the most plausible launching mechanism is based on magnetocentrifugal acceleration in a star-disk system: scenarios include the disk wind \citep{Konigl2000}, X-wind \citep{Shu2000}, and stellar wind \citep{Matt2005, Matt2008a, Matt2008b} models. Jets from a binary system can be explained if the jets emanate from each single star-disk system \citep[e.g. L1551 IRS5:][]{Fridlund1998,Rodriguez1998,Itoh2000,Pyo2002,Pyo2005}, although it has been suggested that one of the binary jets could be destroyed or engulfed by interaction between or merging of the two jets \citep{Murphy2005,Murphy2008}. A single outflow or jet from a binary system could also be produced by a single circumbinary disk \citep{Machida2009,Mundt2010}. \citet{Reipurth2000} has also postulated that the dynamical decay of multiple systems induces outflow activity based on the fact that the binary frequency (79\%-86\%) in 14 giant Herbig-Haro objects is twice higher than the higher multiple frequency. At near-infrared (NIR) wavelengths, \FeII\ and H$_2$ emission lines are excellent tracers of outflows and jets \citep[e.g.][]{Nisini2002,Davis2001,Davis2003,Takami2006,Garcia2008,Garcia2010,Hayashi2009}. Observations obtained with high-angular and high-velocity resolution, particularly those that utilize large ground-base telescopes combined with adaptive optics systems, allow us to study the detailed spatial and kinematical structure close to the launching region \citep{Pyo2003,Pyo2006,Takami2007,Beck2008,Davis2011}. The \HeI\ $\lambda $~1.083 $\mu$m line is a good tracer of the inner hot wind and funnel flow. These are evident as blueshifted absorption combined with emission, and redshifted absorption, respectively \citep{Edwards2003,Dupree2005,Edwards2006,Kwan2007,Fischer2008,Kwan2011}. \citet{Takami2002} found that the \HeI\ emission associated with DG~Tau was spatially extended toward the high-velocity blushifted jet direction. \citet{Pyo2014} will present the detail spatial structure of this \HeI\ extension. Finally, the NIR {\ion{H}{1}} emission lines traces accretion but also outflow activity. For example, spectro-astrometry indicates that the high velocity blueshifted gas observed in Pa$\beta$ emission in DG~Tau is offset along the jet direction \citep{Whelan2004}, while \citet{Beck2010} have detected Br$\gamma$ emission extended by more than 0$\farcs$1 from the star along the jet axes in four CTTSs. In this paper we present NIR 1 $\mu$m spectroscopy over a wavelength range that covers \HeI\ $\lambda$~1.083 $\mu$m, \PaG\ $\lambda$~1.094 $\mu$m, and \FeII\ $\lambda~$1.257 $\mu$m emission lines obtained with the integral field spectrograph, NIFS, at the Gemini North Observatory. We focus on high-resolution \FeII\ emssion maps which show the complicated structure associated with the binary system UY Aur.
\label{sec:discussion} The primary shows deep blueshifted absorption in the \HeI\ $\lambda$~10830 line profile (Figure~\ref{fig_profile}). \citet{Edwards2003,Edwards2006} pointed out that the diversity in width and depth of the blueshifted absorption of \HeI\ indicates an inner wind emanating from a star with a large solid angle. \citet{Kwan2007} classified the \HeI\ line profile they observed towards UY Aur as a disk wind with the second source of emission because the profile in \citet{Edwards2006} showed narrow blueshifted absorption and emission extended to blueward. However, the profile which we observed on 2007 (Figure~\ref{fig_profile}$\textit{a}$) showed more developed blueshifted absorption and resembled the profile of CY~Tau in \citet{Kwan2007}, which were interpreted in terms of a stellar wind model. The \PaG\ $\lambda$~1.094 $\mu$m emission does not show any spatial extension. This is consistent with the interpretation that the Br$\gamma$ and \PaG\ emissions arise from very compact magnetospheric accretion columns in the vicinity of the central stars. On the other hand, \citet{Beck2010} reported that 50\,\% of their targets (four out of eight) had spatially extended Br$\gamma$ emission, suggesting that the atomic Hydrogen emission can originate from extended jets or scattering by extended dust around the central stars. They assumed that the other four objects also had spatially extended Br$\gamma$ emission that was obscured by the bright continuum emission close to the central stars. The extended emission could not have been detected if the physical conditions of the jet gas were not adequate: i.e., if the temperature were too low or the density were too high. In the case of spatially extended dust scattering, the spatial distribution of the line emission should follow that of the continuum emission. The velocity structure between UY Aur A and B is complicated because of the overlap of blueshifted and redshifted emission features. A narrow structure connecting the primary and secondary disks similar to the `bridge' observed in the present studies is predicted by accretion flow simulations of circumbinary disks \citep{Gunther2002,Hanawa2010, Fateeva2011}. The bridge might then be related to the boundary region between two colliding accretion flows. If the circumbinary disk planes are not perpendicular to the line of sight, both blueshifted and redshifted accretion flows would exist around the `bridge' structure. However, it is difficult to produce the high velocity features seen in our data as a result of colliding flows between the two circumstellar disks because the accretion gas flows should basically have Keplerian velocities, of the order of a few \kms\ at the edge of circumstellar disks. Thus, the high velocity emission features are probably associated with outflows launched from the vicinity of the central stars. The \FeII\ emission profile observed towards the UY Aur system exhibits blueshifted emission at $-$154 -- $-$46 \kms\ and redshifted emission at $+$67 -- $+$201 \kms\ (after correcting for the 42$\degree$ inclination angle of the circumbinary disk). Because the terminal velocity of an outflow is comparable to the Keplerian rotation velocity of a disk at the launching region \citep{Kudoh1998}, the fast flow velocities indicate that the launching region is located in a deep potential well, close to the star in the accreting star-disk system. If we use the stellar masses of the primary and secondary as 0.6\,$M_\sun$ and 0.34\,$M_\sun$, respectively, estimated by \citet{HK2003} using the pre-main-sequence tracks of \citet{Siess2000}, and assume three-times larger lever arm radii than the foot point radii for magnetocentrifugal acceleration \citep{Konigl2000}, then the foot points for outflow acceleration correspond to $\sim$~0.22\,--\, 0.6\,AU and $\sim$~0.12\,--\,0.34\,AU for UY Aur A and B, respectively. The other estimates of stellar masses are more or less consistent with that of \citet{HK2003}. Based on the Keplerian rotation measured by the $^{13}$CO gas in the circumbinary disk, \citet{Duvert1998} estimated the total mass of the UY Aur binary as $\sim$~1.2 $M_\sun$. \citet{Close1998} estimated the total mass of the binary as 1.61$^{+0.47}_{-0.67}$ $M_\sun$ using orbital parameters derived from observations obtained since 1944. Together these results show the uncertainty of the total mass estimate of the binary system, not affecting the above estimates of launching radii. The geometrical distribution of \FeII\ emission from the primary suggests a wind with a wide opening angle: the blueshifted and redshifted emission in Figure~\ref{fig_RB} is widely distributed, and in the diagram of $<V>$ (Figure~\ref{fig_Vpeak}$\textit{a}$) the blueshifted emission opens out in a `V' shape around the primary. The gap of 0$\farcs$2 between the primary and the redshifted emission can be explained by the optically thick circumstellar disk of UY Aur A. The FWHM diameter of the star in our images was measured to be $\sim$~0$\farcs$14. We have to be careful when interpreting structures within $\sim$~0$\farcs$1 of the star because the subtraction residuals associated with Poission noise in the central region can hugely affect the structure. However, we can say that the blueshifted gas is spread over more than $\sim$~0$\farcs$2 radius. We note that a similar fast wind with a wide opening angle has been reported from at least one other young source, L1551 IRS 5 \citep{Pyo2005,Pyo2009}. The complicated velocity structure in UY Aur prevents us from easily identifying the origin of the various emission features. The high-velocity redshifted features labelled A and B in Figure~\ref{fig_Vpeak} could be assocaited with the edges of a fast wind with a wide opening angle from UY Aur A. Feature `C' could then be an extension of feature `A'. With this model, however, it is diffcult to explain the blueshifted emission seen between features A, B and C. This emission may be associated with a micro jet from the secondary, UY Aur B. Clearly, deeper observations with spatial and spectral resolutions comparable to those presented here are required in \FeII\ to further probe this complex region. The elongated knot of redshifted emission labelled region `C' in Figure~\ref{fig_Vpeak} may be associated with the primary rather than the secondary, i.e. it may be an extension of region `A'. The secondary may drive only a narrow blueshifted jet towards the North. Deeper observations with spatial and spectral resolutions comparable to the presented studies are required in the \FeII\ line to further probe this complex region. The absence of \FeII\ emission at the south of the secondary (within its Roche lobe) in Figure~\ref{fig_RB}{\em b} is curious. Intrinsically existing emission could have been obscured by a gas-dust cloud ejected from the binary or the secondary. A notably large flux variation of the secondary may be caused buy such a cloud. For example, in 1992 it became 5 mag fainter than the primary in the R-band \citep{Herbst1995}. The R-band magnitude difference ($\Delta$m) between the primary and secondary has changed by 6.6$^m$ in 1996 \citep{Close1998} and 2.0$^m$ in 2000 \citep{Brandeker2003}. In the H-band, the $\Delta$m changed 1.7$^m$ in 1996, 0.43$^m$ in 2002, and 1.7$^m$ in 2005. Overall, the secondary has exhibited larger magnitude variations than the primary: $\pm$1.3$^m$ compared to $\pm$0.3$^m$ between 1996 - 2005 \citep{Hioki2007, Close1998}. These magnitude variations should be related to the changes in mass loss (outflow) and mass accretion rates. \citet{Berdnikov2010} pointed out that a gas-dust cloud had obscured UY Aur during 1945 - 1974. This cloud may still be responsible to obscure UY Aur B. \citet{Skemer2010} reported that the silicate spectrum of UY Aur B is much flatter than expected. This may be the result of foreground extinction for an edge-on disk: the UY Aur disk may be viewed edge-on, although such a disk has not been detected in direct imaging observations of UY Aur B. Thus the disk may not be precisely `edge-on', i.e., the inclination may be larger than $\sim$ 15 degrees with respect to the line of sight. \citet{Monin2007} postulated that the mis-alignment of circumstellar disks in binary systems may be quite common in wide binaries (with a $>$ 100 AU) for Class 0, I , and II phases. UY Aur system may be an example of such a case. Figure~\ref{fig_draw} shows a schematic drawing of the UY Aur system, illustrating our overall understanding of this remarkable object.
14
3
1403.3474
We present high-resolution 1.06-1.28 μm spectra toward the interacting binary UY Aur obtained with GEMINI/NIFS and the adaptive optics system Altair. We have detected [Fe II] λ1.257 μm and He I λ1.083 μm lines from both UY Aur A (the primary source) and UY Aur B (the secondary). In [Fe II] UY Aur A drives fast and widely opening outflows with an opening angle of ~90° along a position angle of ~40°, while UY Aur B is associated with a redshifted knot. The blueshifted and redshifted emissions show a complicated structure between the primary and secondary. The radial velocities of the [Fe II] emission features are similar for UY Aur A and B: ~ -100 km s<SUP>-1</SUP> for the blueshifted emission and ~ +130 km s<SUP>-1</SUP> for the redshifted component. The He I line profile observed toward UY Aur A comprises a central emission feature with deep absorptions at both blueshifted and redshifted velocities. These absorption features may be explained by stellar wind models. The He I line profile of UY Aur B shows only an emission feature. <P />Based on observations obtained at the Gemini Observatory, which is operated by the Association of Universities for Research in Astronomy, Inc., under a cooperative agreement with the NSF on behalf of the Gemini partnership: the National Council (United Kingdom), the National Research Council (Canada), CONICYT (Chile), the Australian Research Council (Australia), CNPq (Brazil), and CONICET (Argentina).
false
[ "UY Aur A", "UY Aur B", "UY Aur", "emission features", "Research", "both UY Aur A", "the interacting binary UY Aur", "Argentina", "stellar wind models", "the National Research Council", "United Kingdom", "~40°", "the Australian Research Council", "Altair", "[Fe II] UY Aur A drives", "Gemini", "deep absorptions", "CONICET", "a central emission feature", "Fe II" ]
9.411354
12.272401
-1
483410
[ "Golkhou, Vahid Z.", "Butler, Nathaniel R." ]
2014ApJ...787...90G
[ "Uncovering the Intrinsic Variability of Gamma-Ray Bursts" ]
57
[ "School of Earth and Space Exploration, Arizona State University, Tempe, AZ 85287, USA; Cosmology Initiative, Arizona State University, Tempe, AZ 85287, USA.;", "Cosmology Initiative, Arizona State University, Tempe, AZ 85287, USA." ]
[ "2014ApJ...794L...8B", "2015A&A...573L...7B", "2015A&A...583A.129Y", "2015Ap&SS.359...20T", "2015ApJ...805...86S", "2015ApJ...811...93G", "2015ApJ...813..127P", "2015JHEAp...7...81G", "2015MNRAS.454.1132T", "2016ApJ...820...68B", "2016ApJ...822...63B", "2016JASS...33..109Y", "2016MNRAS.455L...1W", "2016MNRAS.456.2157I", "2017ApJ...838..143X", "2017ApJ...848L..14G", "2017IJMPD..2630018P", "2018ApJ...857...81G", "2018ApJ...859..163H", "2018MNRAS.475.1708A", "2018MNRAS.478.3525L", "2018PASP..130e1001D", "2019MNRAS.487.5508A", "2019PASP..131f8004A", "2020A&A...636A..82G", "2020ApJ...893...72R", "2020ApJ...893..128A", "2020ApJ...898...14M", "2020ApJS..250....1T", "2020MNRAS.492.4005P", "2021ApJ...918...12F", "2021ExA....52...45C", "2021JInst..16C9007Z", "2021MNRAS.500.5614A", "2021MNRAS.501.4974V", "2021MNRAS.505.4608M", "2022Galax..10...78S", "2022Galax..10...93S", "2022JApA...43...37I", "2022JCAP...02..010S", "2022MNRAS.511..964X", "2022MNRAS.512.2337D", "2022MNRAS.514.2397X", "2022NatCo..13.5611D", "2022Natur.612..228T", "2022arXiv220611088S", "2023A&A...671A.112C", "2023A&A...677A..19J", "2023ApJ...954L..29D", "2023ApJ...956...42S", "2023ApJS..268....5X", "2023Univ....9..245T", "2024ApJ...967..128X", "2024ApJS..270....3X", "2024MNRAS.52711915X", "2024MNRAS.528.1388X", "2024PASJ...76..550O" ]
[ "astronomy" ]
9
[ "gamma-ray burst: general", "methods: data analysis", "methods: statistical", "Astrophysics - High Energy Astrophysical Phenomena", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
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[ "10.1088/0004-637X/787/1/90", "10.48550/arXiv.1403.4254" ]
1403
1403.4254_arXiv.txt
\label{sec:intro} Gamma-Ray Burst (GRB) light curves show a remarkable morphological diversity. While a significant number of bright long bursts ($\sim15\%$) exhibit a single smooth pulse structure, in most cases GRBs appear to be the result of a complex, seemingly random distribution of several pulses. Burst pulses are commonly described as having fast-rise exponential-decay (FRED) shapes \citep[e.g.,][]{1996ApJ...473..998F}. Parameterized analyses of pulse profiles have shown broad log-normal distributions among different bursts and even within a single burst \citep[see, e.g.,][]{1996ApJ...459..393N,2012ApJ...744..141B}. Several approaches have been utilized to characterize the distribution of power versus timescale for GRBs and other astrophysical sources. These include structure function (SF) analyses \citep{1994ApJ...433..494T,1994MNRAS.268..305H,1996A&A...306..395C,1997MNRAS.286..271A}, autocorrelation function (ACF) analyses \citep{1993ApJ...408L..81L,1995ApJ...448L.101F,1996ApJ...464..622I,2004A&A...418..487B,2012ApJ...749..191C}, and Fourier power spectral density (PSD) analyses \citep{2000ApJ...535..158B,2001ApJ...557L..85C,2010ApJ...722..520A,2012MNRAS.422.1785G,2013MNRAS.431.3608D}. In principle, the ACF contains the same information as the PSD, since one is the Fourier transform of the other \citep[the Wiener-Khinchin theorem,][]{chatfield2003analysis}. The SF is mathematically very similar to the ACF. As \citet{2002MNRAS.329...76H} summarizes, the SF, ACF, and PSD -- when calculated for a given dataset -- are not completely equivalent because of time-windowing effects and the presence of measurement noise. For long runs of evenly spaced data, the PSD is used in preference to the ACF, as it can be easier to interpret and understand errors. In cases of short or inhomogeneous data sets, the ACF can provide a more stable measurement. However, as ACF values at different time lags are not statistically independent of each other, the ACF interpretation may not be simple. The first-order SF was introduced in astronomy by \citet{1985ApJ...296...46S}. It has been widely used in the analysis of quasar light curves \citep[e.g.,][]{1994ApJ...433..494T,2002MNRAS.329...76H} and microlensing statistics \citep[e.g.,][]{2001MNRAS.320...21W}. Compared to power-spectral analyses, the SF approach is less dependent on the time sampling \citep{1999ASPC..159..293P}. Following these studies, we define the first-order SF as a measure of the mean square difference of a signal $X(t)$ on timescale (or ``lag'') $\tau$: \be {\rm SF}(\tau)=\left\langle [X(t)-X(t+\tau)]^2 \right\rangle \label{Eq:sf} \ee Here, $\left\langle . \right\rangle$ denotes an averaging over $t$. In Figure \ref{fig:sfp}, we reproduce the typical shape of an SF, from \citet{1992ApJ...396..469H}. \begin{figure} \begin{center} \includegraphics[width=.3\textwidth]{schematic_sf.pdf} \caption{\small Schematic showing a typical SF for a time-series, from \citet{1992ApJ...396..469H}. At short lag-times, the SF flattens due to the measurements error. At long lag-times, the SF again flattens out at a level corresponding to the total variance in the signal. Between these lag-times, the slope of the SF depends on the noise properties of the signal and can be used to identify timescales of interest.} \label{fig:sfp} \end{center} \end{figure} We will be primarily interested below in using the SF to infer the shortest timescale at which a GRB exhibits {\it uncorrelated temporal variability}. In a seminal study, \citet{2000ApJ...537..264W} \citep[and more recently,][]{2013MNRAS.432..857M} utilize Haar wavelet scaleograms to measure minimum timescales. Wavelets are a set of mathematical functions, which form an orthonormal basis to compactly describe narrow time features \citep[e.g.,][]{1992tlw..conf.....D,1994ApJ...424..540N,1997ApJ...483..340K,2013ApJ...764..167S}. Making the connection between the Haar wavelet scaleogram and the SF, as we do below in mathematical detail, sheds new light on prior work, allowing for a more rigorous analysis and better physical interpretation of the signal power versus timescale. We also exploit the large sample of \textit{Swift} GRBs with measured redshifts to perform this analysis, for the first time, in the GRB source frame. A general feature we observe in our scaleograms, provided there is sufficient signal-to-noise ratio (SNR), is a linear rise phase relative to the Poisson noise floor on the shortest timescales (see, e.g., Figures \ref{fig:pick_burst} and \ref{fig:simburst_30s}). We take this to indicate a typical smoothness on the shortest observed timescales. We thus make an essential distinction -- not made in prior studies -- between correlated variability (i.e., smooth or continuous) and uncorrelated variability (e.g., pulses or changes in sign). For example, an exponentially decaying GRB light curve pulse with a fairly long time constant (say 100 s) will still exhibit power (i.e., yield a non-zero SF) on much shorter timescales (say 1 s), provided the SNR is sufficiently large for this to be measured. In contrast, the meaningful timescale (in this case $\approx$ 100 s and not 1 s) is the shortest timescale at which the signal becomes uncorrelated. A simple Taylor expansion of the SF assuming a temporally-smooth signal $X(t)$, shown in Equation \ref{Eq:taylor_expn}, elucidates how the minimum timescale for uncorrelated variability is connected to the scaleogram linear-rise phase. \be X(t+\tau) = X(t) + \tau X'(t)|_{\tau} + ..., \label{Eq:taylor_expn} \ee Substituting Equation \ref{Eq:taylor_expn} into Equation \ref{Eq:sf} and ignoring higher order terms produces Equation \ref{Eq:prop}: \be \sqrt{{\rm SF}(\tau)} \propto \tau \label{Eq:prop} \ee which shows that for timescales where the signal is smoothly varying, we expect a linear dependence on the time lag $\tau$. When the variation becomes non-smooth, SF flattens, providing a signature of the true GRB minimum timescale. Previous studies \citep{2000ApJ...537..264W,2013MNRAS.432..857M} -- which overlook the importance of the ${\rm SF} \propto \tau$ region -- incorrectly interpret the GRB minimum as the shortest timescale at which the SF is first non-zero (after subtracting the measurement noise level), potentially under-estimating the true variability timescales. In this paper, we begin with a more detailed description of our method -- the Haar wavelet structure function -- which exploits a non-decimated, discrete Haar wavelet transform to estimate the SF and, in turn, the minimum variability timescale for a large number GRBs. We discuss the robustness of the structure function in extracting this timescale even in the case of complex GRBs containing multiple, overlapping pulses and we demonstrate self-consistency as the SNR is varied. Next, we apply the methodology to the full sample of GRBs observed by \textit{Swift} BAT, summarizing the derived timescales for the population in the observer and GRB source frames. We conclude by discussing how these minimum variability timescales can elucidate the GRB central engine, help constrain models for the emission mechanism, and potentially also enable a measurement of cosmological time-dilation. \begin{figure} \begin{center} \includegraphics[width=.49\textwidth]{Haar_step.png} \caption{\small A schematic representation of relation between the Haar wavelet coefficients (Equation \ref{Eq:coef}) and the first-order structure function. The Haar mother wavelet as shown step function (with different style) operates using scaling and dilation on a time-series: $\{X_i | i=0..8\}$. $S_i = \sum_{0}^{i}{X_i}$.} \label{fig:ep_pred2} \end{center} \end{figure}
Using a technique based on Haar wavelets, we have studied the temporal properties of a sample of GRB hard X-ray, prompt-emission light curves captured by the BAT instrument on {\it Swift} prior to October 27, 2013. Our approach averages over the time-series captured for a given GRB, providing robust measures of minimum variability timescales. In contrast to previous studies \citep{2000ApJ...537..264W,2013MNRAS.432..857M,2013arXiv1307.7618B}, which simply define the minimum timescale in reference to the measurement noise floor, our approach identifies the signature of temporally-smooth features in the wavelet scaleogram and then additionally identifies a break in the scaleogram on longer timescales as signature of a true, temporally-unsmooth light curve feature or features. We find that this timescale ($\Delta t_{\rm min}$) tends to correspond to the rise-time of the narrowest GRB pulse \citep[see, also,][]{2013arXiv1307.7618B}. We find a median minimum timescale for long-duration GRBs in the source (observer) frame of $\Delta t_{\rm min}=0.5$ s ($\Delta t_{\rm min}=2.5$ s). A consistent value in the source-frame for short-duration GRBs may indicate a common central engine. We find that very few -- at most 15\% (5\%) in the observer frame (source frame) -- of {\it Swift}~GRBs can have minimum timescales below 10 ms. Our timescales are thus considerably longer than the millisecond variability timescales found by \citet{2000ApJ...537..264W} to be common in bright BATSE GRBs. Partial explanation for this discrepancy must come from the fact that {\it Swift}~BAT operates in a lower photon energy range than BATSE, and GRB pulses are known to be more narrow in higher energy bandpasses \citep[e.g.,][]{1996ApJ...459..393N,1995ApJ...448L.101F,1996ApJ...473..998F}. Nonetheless, we note that the variability found in \citet{2000ApJ...537..264W} is not linked to the presence of discernible features in a given light curve (e.g., the pulse rise-time, which it is actually stated to be considerably less than). Given our new distinction between a \citet{2000ApJ...537..264W} type timescale (which we call $\Delta t_{\rm snr}$) -- the minimum possible observable $\Delta t_{\rm min}$ for a GRB of given brightness and not necessarily the true $\Delta t_{\rm min}$ -- it is natural to expect that \citet{2000ApJ...537..264W} have underestimated their minimum timescales. We note that our minimum timescales are broadly consistent with those found in pulse-fitting studies \citep[e.g.,][]{1995ApJ...448L.101F,1996ApJ...459..393N}. \subsection{Constraints on the Fireball Model} The standard fireball model postulates the release of a large amount of energy by a central engine into a concentrated volume \citep{1978MNRAS.183..359C,2004RvMP...76.1143P}, which causes the resulting outflow to expand and quickly become relativistic \citep{1986ApJ...308L..43P}. These relativistic expanding shells -- with different Lorentz factors -- in general collide, resulting in Gamma-ray flares and potentially rich temporal structure. Our extracted timescales ($\Delta t_{\rm min}$) should provide a diagnostic on the central engine power and its evolution. The size of the central engine is limited to $R < c \, \Delta t_{\rm min}$, which for the smallest minimum variability timescale derived above \citep[$\sim 10$ ms; see also,][]{2002MNRAS.330..920N} is $R < 3 \times 10^3$ km. Typical $\Delta t_{\rm min}$ values from above lead to $R < 2 \times 10^5$ km. For the first time, due to a large sample of GRBs with measured redshift, we are able to perform these calculations in the source frame. In the source frame, we are not able to confirm that the minimum variability timescale of short-duration GRBs is substantially shorter than that of long-duration. Hence, we cannot demonstrate that short-duration GRBs have a more condensed central engine than the former \citep[see,][]{2013arXiv1307.7618B}. We can derive additional constraints on the GRB emission region, following the discussion in \citet{2000ApJ...537..264W}. In the ``external shock'' picture, shells of material produced by the GRB impact material in the external medium. The physical dimension of clouds and their patchiness -- in the direction perpendicular to the expansion of the shell -- is constrained by $\Delta t_{\rm min}$. If we assume a single shell expanding at very close to light speed, the arrival time for photons from the shell will be calculated as the summation of travel time of the shell to the radius of impact and the travel time of the Gamma-rays to Earth. Photons from off-axis regions of the relativistic expanding shell experience a purely geometrical delay compared with photons from on-axis regions \citep{2004RvMP...76.1143P} reaching the observer. The observed delay depends only on the radius of the shell $R$ at the time of impact with a cloud in the external medium and the angular radius of the Gamma-ray emission region as subtended from the burst site ($\Delta \Theta$). \citet{2000ApJ...537..264W} report millisecond variability superposed on pulses of significantly longer rise-times. High cloud patchiness can potentially explain this modulation, implying $\Delta \Theta < \Delta t_{\rm min}/(2\Gamma\, T_{\rm rise}) < 0.0002$ radians, where $\Gamma$ is the bulk Lorentz factor. Only $\sim 5 \times 10^{-3}$ of the emitting shell is active at a given time \citep[also,][]{1999ApJ...512..683F}. However, for the bursts in our sample, with typical minimum variability timescale $\sim 0.1 \, s$, $T_{\rm rise} > 1 \, s$, and assuming $\Gamma > 100$, we find that $\Delta \Theta < 5 \times 10^{-3}$ radians which is comparable with the typical surface filling factor \citep{1999ApJ...512..683F}. In terms of the external shock scenario, the extracted $\Delta t_{\rm min}$ can circumscribe the size scale of the impacted cloud along the line of sight. For a thin shell, the Gamma-ray radiation will start when the relativistic shell hits the inner boundary of the cloud with the peak flux produced as the shell reaches the densest region or center of the cloud. The size scale of the impacted cloud is limited by $2 \Gamma^2 c\,\Delta t_{\rm min}$ since the shock is moving near light speed \citep{1996ApJ...473..998F}. For the smallest $\Delta t_{\rm min}$ found $\sim 10$ ms, and assuming $\Gamma < 1000$, the cloud size must be smaller than $40$ AU. In the ``internal shock'' scenario \citep[e.g.,][]{1994ApJ...430L..93R}, the relativistic expanding outflow released from a central engine is assumed to be variable, consisting of multiple shells of differing $\Gamma$. These shells propagate and expand adiabatically until a faster shell collides with a slower one, resulting in a measurable rise time. This rise time to an outside observer would appear as: $\Delta t_{r1} \approx \Delta R / 2c\Gamma_1^2$, where $\Delta R$ and $\Gamma_1$ are the thickness and resulting Lorentz factor of the merged shell \citep[e.g.,][]{2007ApJ...667.1024K}. Assuming the same scenario but for two other shells yields $\Delta t_{r2} \approx \Delta R / 2c\Gamma_2^2$. Writing $\Delta \Gamma = \Gamma_1 - \Gamma_2$, we have $\Delta \Gamma / \Gamma \approx 1/2 \, (\Delta t_{\rm min} / T_{\rm rise})$. In \citet{2000ApJ...537..264W}, the ratios $\Delta t_{\rm min} / T_{\rm rise}$ were argued to be small, implying a narrow dispersion in $\Gamma$. We, however, find $\Delta t_{\rm min} \sim T_{\rm rise}$, suggesting instead a broad range of possible Lorentz factors. Finally, we find evidence that our minimum timescales correlate with redshift, possibly providing indication of cosmological time-dilation. However, the measurement threshold also appears to correlate strongly with redshift. This indicates that threshold effects likely dominate the apparent correlation and that the correlation may not be real. It is possible that additional features present in the Haar scaleogram (slopes, breaks on longer timescales) -- which richly describe the full GRB light curve and not just the minimum timescale -- may yield correlations with intrinsic quantities like redshift. We will study this further in future work.
14
3
1403.4254
We develop a robust technique to determine the minimum variability timescale for gamma-ray burst (GRB) light curves, utilizing Haar wavelets. Our approach averages over the data for a given GRB, providing an aggregate measure of signal variation while also retaining sensitivity to narrow pulses within complicated time series. In contrast to previous studies using wavelets, which simply define the minimum timescale in reference to the measurement noise floor, our approach identifies the signature of temporally smooth features in the wavelet scaleogram and then additionally identifies a break in the scaleogram on longer timescales as a signature of a true, temporally unsmooth light curve feature or features. We apply our technique to the large sample of Swift GRB gamma-ray light curves and for the first time—due to the presence of a large number of GRBs with measured redshift—determine the distribution of minimum variability timescales in the source frame. We find a median minimum timescale for long-duration GRBs in the source frame of Δt <SUB>min</SUB> = 0.5 s, with the shortest timescale found being on the order of 10 ms. This short timescale suggests a compact central engine (3 × 10<SUP>3</SUP> km). We discuss further implications for the GRB fireball model and present a tantalizing correlation between the minimum timescale and redshift, which may in part be due to cosmological time dilation.
false
[ "minimum variability timescales", "longer timescales", "complicated time series", "cosmological time dilation", "Haar wavelets", "wavelets", "features", "the minimum variability timescale", "measured redshift", "a median minimum timescale", "the minimum timescale", "GRB", "Swift GRB gamma-ray light curves", "narrow pulses", "GRBs", "This short timescale", "the shortest timescale", "redshift", "signal variation", "the wavelet scaleogram" ]
16.138189
1.052724
5
482987
[ "Yuan, Qiang", "Huang, Xiaoyuan", "Liu, Siming", "Zhang, Bing" ]
2014ApJ...785L..22Y
[ "Fermi Large Area Telescope Detection of Supernova Remnant RCW 86" ]
36
[ "Key Laboratory of Particle Astrophysics, Institute of High Energy Physics, Chinese Academy of Sciences, Beijing 100049, China; Department of Physics and Astronomy, University of Nevada, Las Vegas, NV 89154, USA", "Key Laboratory of Dark Matter and Space Astronomy, Purple Mountain Observatory, Chinese Academy of Sciences, Nanjing 210008, China", "Key Laboratory of Dark Matter and Space Astronomy, Purple Mountain Observatory, Chinese Academy of Sciences, Nanjing 210008, China", "Department of Physics and Astronomy, University of Nevada, Las Vegas, NV 89154, USA" ]
[ "2014MNRAS.444..860A", "2014NuPhS.256...56M", "2014NuPhS.256..186B", "2014arXiv1404.4924K", "2015A&A...573A..37Y", "2015A&A...580A..74A", "2015APh....65...80M", "2015ARNPS..65..245F", "2015ApJ...809..102L", "2015ApJ...813....3A", "2015CRPhy..16..674H", "2015JPhCS.642a2018M", "2016ApJ...817...64X", "2016ApJ...818...63B", "2016ApJ...819...98A", "2016ApJ...820L...3Y", "2016ApJ...826...63L", "2017AIPC.1792h0007B", "2017ApJ...835...42G", "2017ApJ...851..100C", "2017JHEAp..15....1S", "2018A&A...612A...4H", "2018ApJ...853....2G", "2018ApJ...867...55X", "2018AsBio..18..779W", "2018NPPP..297..158B", "2019ApJ...876...24Z", "2019ApJ...887...47C", "2020PASJ...72...72S", "2021ApJ...917...96S", "2022ApJ...924...45S", "2022ApJ...938...59S", "2022ApJ...941..194X", "2022arXiv220209530W", "2023PASJ...75.1344B", "2024Univ...10..203G" ]
[ "astronomy" ]
5
[ "cosmic rays", "gamma rays: ISM", "ISM: supernova remnants", "radiation mechanisms: non-thermal", "Astrophysics - High Energy Astrophysical Phenomena" ]
[ "1973ApJS...26...19V", "1975AuJPA..37...39C", "1987A&A...183..118K", "1996A&A...315..243R", "1996A&AS..118..329W", "1997A&A...328..628V", "1997AJ....114.2664S", "2000A&A...360..671B", "2000PASJ...52.1157B", "2001ApJ...546..447D", "2001ApJ...550..334B", "2001RPPh...64..429M", "2002APh....16..429B", "2002ApJ...581.1116R", "2002ISAA....5.....S", "2003A&A...407..249S", "2003MNRAS.342.1117M", "2006ChJAA...6..635Z", "2008PASJ...60S.123Y", "2009ApJ...692.1500A", "2009ApJ...697.1071A", "2009Sci...325..719H", "2010ApJ...710L.151T", "2010ApJ...712..287E", "2010ApJ...717..372C", "2011ApJ...731...32F", "2011ApJ...734...28A", "2011ApJ...735..120Y", "2011ApJ...740L..12W", "2011ApJ...740L..51T", "2011ApJ...741...96W", "2012A&A...545A..28L", "2012ApJ...744...80A", "2012ApJ...744L...2G", "2012ApJ...752..135K", "2012ApJ...759...89H", "2012ApJ...761..133Y", "2012ApJS..199...31N", "2012MNRAS.425.2810A", "2013A&A...553A..34D", "2013ApJ...774...36C", "2013ApJ...779..179P", "2013FrPhy...8..714R", "2013MNRAS.434.2202A", "2013Sci...339..807A", "2014A&A...562A.141M", "2014ApJ...781...64X", "2014ApJ...783...32A" ]
[ "10.1088/2041-8205/785/2/L22", "10.48550/arXiv.1403.4915" ]
1403
1403.4915_arXiv.txt
Supernova remnants (SNRs) are believed to be the most probable candidates of the Galactic cosmic ray (CR) acceleration sources. However, direct observational evidence is not available until there are $\gamma$-ray detections of SNRs \citep[e.g.,][]{2010ApJ...710L.151T,2013Sci...339..807A}. Up to now, nearly $20$ SNRs have been discovered in TeV $\gamma$-ray band, among which $7$ are firmly identified as shell-type SNRs \citep{2013FrPhy...8..714R} and about half are interacting with molecular clouds\footnote{http://tevcat.uchicago.edu}. In the two-year catalog of Fermi-LAT (2FGL), there are $6$ firmly identified SNRs based on the spatial extension and $4$ associated point-like SNRs \citep{2012ApJS..199...31N}. Additionally there are $59$ 2FGL sources which might be associated with SNRs based on the spatial match between the error circles of 2FGL sources and the SNR extensions \citep{2012ApJS..199...31N}. With the accumulation of Fermi-LAT data, more and more SNRs were detected \citep{2010ApJ...717..372C, 2011ApJ...740L..51T,2011ApJ...734...28A,2011ApJ...740L..12W, 2012ApJ...744L...2G,2012ApJ...744...80A,2012ApJ...752..135K, 2012ApJ...759...89H,2013ApJ...774...36C,2013MNRAS.434.2202A, 2013ApJ...779..179P,2014ApJ...781...64X,2014ApJ...783...32A}. Although the $\gamma$-ray emission mechanism of individual SNRs is subject to debate, it is possible to approach the nature of $\gamma$-ray emission of SNRs through a population study with a large sample of $\gamma$-ray SNRs \citep{2012ApJ...761..133Y,2013A&A...553A..34D}. Increasing the sample of $\gamma$-ray SNRs can be essential for understanding their non-thermal characteristics. The shell-type SNR G315.4-2.3, also known as RCW 86, is a young remnant probably associated with supernova SN 185 \citep{2002ISAA....5.....S, 2006ChJAA...6..635Z}. The angular diameter of this SNR is about $42'$, with a clear shell in radio \citep{1987A&A...183..118K,1996A&AS..118..329W,2001ApJ...546..447D}, infrared \citep{2011ApJ...741...96W}, optical \citep{1973ApJS...26...19V, 1997AJ....114.2664S} and X-ray bands \citep{1997A&A...328..628V, 2000A&A...360..671B,2000PASJ...52.1157B,2001ApJ...550..334B, 2002ApJ...581.1116R}. The distance of RCW 86 is estimated to be $2.3-2.8$ kpc throught optical spectroscopy observations \citep{1996A&A...315..243R,2003A&A...407..249S}. In the very high energy (VHE) $\gamma$-ray band, a well extended source with morphology consistent with the X-ray image has been revealed by HESS \citep{2009ApJ...692.1500A}. The spectral index of VHE $\gamma$-rays is about $2.5$ and the flux is about $10\%$ of that of the Crab nebula \citep{2009ApJ...692.1500A}. \citet{2012A&A...545A..28L} analyzed $\sim3$ year Fermi-LAT data and found no significant excess from this SNR. Upper limits of $\gamma$-ray flux in the GeV band were derived \citep{2012A&A...545A..28L}. With multi-wavelength observations, the high energy radiation mechanism and particle acceleration can be studied \citep{2009ApJ...692.1500A, 2012A&A...545A..28L}. Here we report the detection of GeV $\gamma$-ray emission from RCW 86, with $5.4$ year Fermi-LAT data. The data analysis, including the morphology and the spectrum, is presented in Sec. 2. Based on the $\gamma$-ray spectrum and the multi-wavelength spectral energy distribution (SED) of RCW 86, we discuss its non-thermal emission mechanism in Sec. 3. Finally Sec. 4 is the conclusion.
In this work we report the detection of GeV $\gamma$-rays from a shell-type SNR RCW 86 with Fermi-LAT. Analyzing $5.4$ year Fermi-LAT data, we find an extended source coincident with the radio or VHE $\gamma$-ray image of RCW 86 with a significance higher than $5\sigma$. The point source assumption is less favored than the extended source assumption. The GeV $\gamma$-ray spectrum is found to be very hard, $\Gamma\approx1.4\pm0.2$. The multi-wavelength SED from radio to VHE $\gamma$-rays can be well described with a simple one zone leptonic model. The hadronic scenario may face difficulty in producing the very hard GeV $\gamma$-ray spectrum and in the total energy budget given that the environmental medium density is relatively low. The analogy of the non-thermal GeV-TeV spectrum and the lack of strong thermal X-ray emission of this SNR with several other shell-type SNRs makes them form a distinct class of $\gamma$-ray SNRs, which may point to the nature of the particle acceleration and radiation in SNRs.
14
3
1403.4915
Using 5.4 yr Fermi Large Area Telescope data, we report the detection of GeV γ-ray emission from the shell-type supernova remnant RCW 86 (G315.4-2.3) with a significance of ~5.1σ. The data slightly favors an extended emission of this supernova remnant. The spectral index of RCW 86 is found to be very hard, Γ ~ 1.4, in the 0.4-300 GeV range. A one-zone leptonic model can well fit the multi-wavelength data from radio to very high energy γ-rays. The very hard GeV γ-ray spectrum and the inferred low gas density seem to disfavor a hadronic origin for the γ-rays. The γ-ray behavior of RCW 86 is very similar to several other TeV shell-type supernova remnants, e.g., RX J1713.7-3946, RX J0852.0-4622, SN 1006, and HESS J1731-347.
false
[ "HESS J1731", "very high energy γ-rays", "several other TeV shell-type supernova remnants", "the γ-rays", "RX", "RCW", "~5.1σ", "GeV γ-ray emission", "γ-ray spectrum", "Fermi Large Area Telescope data", "the shell-type supernova remnant RCW", "HESS", "Fermi Large Area Telescope", "SN", "the inferred low gas density", "G315.4", "TeV", "this supernova remnant", "The γ-ray behavior", "radio" ]
4.382933
4.773156
60
524399
[ "Sahni, Varun" ]
2014arXiv1403.1537S
[ "Ya. B. Zeldovich (1914-1987): Chemist, Nuclear Physicist, Cosmologist" ]
0
[ "-" ]
null
[ "astronomy", "physics" ]
4
[ "Physics - History and Philosophy of Physics", "Astrophysics - Cosmology and Extragalactic Astrophysics", "General Relativity and Quantum Cosmology", "High Energy Physics - Phenomenology" ]
[ "1972CoASP...4..173S", "1988Natur.331..671S", "1989RvMP...61..185S", "1995PhR...262....1S", "2005ApJ...624..463G", "2009A&A...500...19S" ]
[ "10.48550/arXiv.1403.1537" ]
1403
1403.1537_arXiv.txt
Yakov Borisovich Zeldovich was very talented. His active scientific career included major contributions in fields as diverse as chemical physics (adsorption \& catalysis), the theory of shock waves, thermal explosions, the theory of flame propogation, the theory of combustion \& detonation, nuclear \& particle physics, and, during the latter part of his life: gravitation, astrophysics and cosmology \cite{zel}. \z made key contributions in all these area's, nurturing a creative and thriving scientific community in the process. His total scientific output exceeds 500 research article and 20 books. Indeed, after meeting him, the famous English physicist Stephen Hawking wrote ``Now I know that you are a real person and not a group of scientists like the Bourbaki''.$^1$\footnotetext[1]{Bourbaki was the pseudonym collectively adopted by a group of twentieth century mathematicians who wrote several influential books under this pseudonym on advanced mathematical concepts.} Others have compared his enormously varied scientific output to that of Lord Raleigh who, a hundred years before \z, worked on fields as varied as optics and engineering. Remarkably \z never received any formal university education ! He graduated from high school in St. Petersburg at the age of 15 after which he joined the {\em Institute for Mechanical Processing of Useful Minerals} (`Mekhanabor') to train as a laboratory assistant. The depth of \z's questioning and his deep interest in science soon reached senior members of the scientific community and, in 1931, the influential soviet scientist A.F. Ioffe wrote a letter to Mekhanabor requesting that \z be ``released to science''. \z defended his PhD in 1936 and, years later, reminiscenced of the ``happy times when permission to defend [a PhD] was granted to people who had no higher education''. Despite his never having been formally taught (or perhaps because of it !) \z developed a very original style of doing science, and became, in the process, an exceptional teacher. It is also interesting that in his early years \z had been an experimentalist as well as a theoretician, and this closeness to complementary aspects of science guided him throughout his later life.
14
3
1403.1537
Ya.B. Zeldovich was a pre-eminent Soviet physicist whose seminal contributions spanned many fields ranging from physical chemistry to nuclear and particle physics, and finally astrophysics and cosmology. March 8, 2014 marks Zeldovich's birth centenary, and this article attempts to convey the zest with which Zeldovich did science, and the important role he played in fostering and mentoring a whole generation of talented Scientists.
false
[ "talented Scientists", "Scientists", "physical chemistry", "many fields", "B. Zeldovich", "Zeldovich", "nuclear and particle physics", "science", "a pre-eminent Soviet physicist", "Soviet", "Zeldovichs birth centenary", "a whole generation", "the important role", "whose seminal contributions", "March", "the zest", "this article", "2014 marks", "March 8, 2014 marks", "he" ]
11.536507
2.624293
-1
800355
[ "Bezrukov, F. L.", "Gorbunov, D. S." ]
2014PhLB..736..494B
[ "Relic gravity waves and 7 keV dark matter from a GeV scale inflaton" ]
62
[ "CERN, CH-1211 Genève 23, Switzerland", "Institute for Nuclear Research of the Russian Academy of Sciences, Moscow 117312, Russia" ]
[ "2014JCAP...08..040D", "2014JCAP...10..013C", "2014JCAP...11..033C", "2014JHEP...08..086G", "2014JHEP...09..006C", "2014PhLB..738..118L", "2014PhLB..738..380F", "2014PhRvD..89l1302C", "2014PhRvD..89l7305A", "2014PhRvD..90a1701B", "2014PhRvD..90c5002D", "2014PhRvD..90c5005R", "2014PhRvD..90c5009C", "2014PhRvD..90c5014D", "2014arXiv1405.3730B", "2014arXiv1409.3656P", "2015EPJC...75...31F", "2015JCAP...05..053C", "2015JCAP...09..066P", "2015JHEP...01..006A", "2015JHEP...03..064M", "2015JHEP...08..023B", "2015JHEP...08..072A", "2015PhLB..744..218L", "2015PhLB..745...29G", "2015PhLB..749..283M", "2015PhRvD..91k5013P", "2015PhRvD..92g5005B", "2015PhRvD..92i5018R", "2015PhRvD..92j3509P", "2015PhRvL.115p1802A", "2015arXiv151005483C", "2015arXiv151100281M", "2016AnP...528..161R", "2016JCAP...04..003M", "2016JCAP...11..038K", "2016JPhCS.770a2032L", "2016PhRvD..94l5022H", "2016RPPh...79l4201A", "2016chep.confE.826M", "2017ApJ...836..217C", "2017JCAP...01..025A", "2017JHEP...06..050F", "2017JInst..12P5025C", "2017JPhCS.800a2018P", "2017JPhCS.873a2019P", "2017PhLB..769...67B", "2017PhRvD..95g1101A", "2017PhRvD..96c5018H", "2017PhRvD..96k5019B", "2018JCAP...02..002B", "2018JCAP...04..006B", "2018PhRvL.120f1301B", "2018SCPMA..61j1005G", "2019PrPNP.104....1B", "2019arXiv190310712P", "2020PhRvD.101c5009J", "2020PhRvD.102k5011B", "2021PhLB..82036524G", "2021PhRvD.103f3512B", "2023JCAP...02..055B", "2023PhLB..84638206L" ]
[ "astronomy", "physics" ]
4
[ "High Energy Physics - Phenomenology", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
[ "1985PhLB..157..458B", "1999PhRvL..82.2832S", "2001ApJ...562..593A", "2004PhRvD..69l3523T", "2005PhLB..620...17A", "2006PhLB..638..401A", "2006PhLB..639..414S", "2007JHEP...01..091A", "2009ARNPS..59..191B", "2009JCAP...05..012B", "2009PhLB..671..211A", "2010JHEP...05..010B", "2012PhLB..716....1A", "2012PhLB..716...30C", "2012PhRvD..86a0001B", "2013JHEP...07..140B", "2013PhRvD..88d3502V", "2014A&A...571A..22P", "2014ApJ...789...13B", "2014PhRvL.112x1101B", "2014PhRvL.113y1301B" ]
[ "10.1016/j.physletb.2014.07.060", "10.48550/arXiv.1403.4638" ]
1403
1403.4638_arXiv.txt
Discovery of the neutral scalar with properties very close to what we expect for the SM Higgs boson \cite{Aad:2012tfa,Chatrchyan:2012ufa} and absence of any definite hints of supersymmetry at LHC asks for its replacement as a solution to gauge hierarchy problem. Some hope is associated with conformal or scale invariance that might be a symmetry of the SM at tree level, but for the only dimensionfull parameter of the SM $v$ which gives the vacuum expectation value to the Englert--Brout--Higgs (EBH) field. Yet this parameter may be generated by the vacuum expectation value of a new scalar field (scale invariance breaking messenger) introduced into particle physics, so that the SM sector is scale invariant at tree level. The field itself may be used to solve other SM problems. Here we discuss the idea that it may serve as an inflaton in the early Universe. The renormalizable model realizing this idea was suggested in \cite{Shaposhnikov:2006xi} and further developed in \cite{Anisimov:2008qs,Bezrukov:2009yw,Bezrukov:2013fca}. In this Letter we assume that the SM sector of the model is scale invariant, in order to alleviate the hierarchy problem of the Higgs mass. The only violation of scale invariance is assumed to be present in the inflaton, which can be considered as a messenger of the scale symmetry breaking, exact mechanism of this breaking is beyond the present analysis. Technically this means that we assumed that other dimensionful parameters, like Higgs boson mass term or cubic terms in the potential are small and can be neglected. Introducing these terms at electroweak scale would not change the phenomenology of the model. With only one dimensionfull parameter explicitly breaking scale invariance in the inflaton sector, the model is consistent with cosmological observations \cite{Bezrukov:2013fca} and constraints from particle physics related to the possible manifestation of the light inflaton in B-meson decays \cite{Bezrukov:2009yw}. The model may be further extended by introducing three Majorana fermions $N_I$, $I=1,2,3$, which are singlets with respect to the SM gauge group. Yukawa-type coupling to inflaton provides these fermions with Majorana mass terms when the inflaton field obtains vacuum expectation value. The Yukawa-type couplings between $N_I$, EBH doublet, and SM lepton doublets lead to Dirac masses for neutrinos, and the active neutrino masses are then obtained from seesaw type I formula \cite{Minkowski:1977sc}. Hence the fermions serve as sterile neutrinos, and the Yukawa couplings in a part of the parameter space may explain the baryon asymmetry of the Universe via leptogenesis, e.g.\ implementing the $\nu$MSM scheme \cite{Asaka:2005pn,Boyarsky:2009ix}. Remarkably, the lightest sterile neutrino $N_1$, provided tiny coupling to active neutrinos, may serve as non-thermal DM produced by inflaton decays in the early Universe \cite{Shaposhnikov:2006xi}. Therefore the suggested model with seven new degrees of freedom added to the SM explains the neutrino oscillations, DM phenomena, baryon asymmetry of the Universe and exhibits the inflationary dynamics at early times thus solving the Hot Big Bang theory problems. Given the allowed range of the inflaton mass the DM sterile neutrino is naturally light here, $1\keV<M_1<1\MeV$. In this \emph{Letter} we discuss the particular choice of $M_1=7\keV$ motivated by recently found anomalous line in cosmic X-ray spectra of galaxy clusters and Andromeda galaxy observed by orbital telescopes \cite{Abazajian:2001vt,Bulbul:2014sua,Boyarsky:2014jta}. We outline the viable region of the model parameter space consistent with this choice of sterile neutrino mass and give definite predictions for the inflationary cosmological parameters and the inflaton mass, its lifetime and branching ratio of B-meson to inflaton which {\em allow to thoroughly investigate this model}. Remarkably, recent results of BICEP2 experiment \cite{Ade:2014xna} on detection of B-mode polarization, interpreted as primordial tensor perturbations, can completely fix all the parameters of the model. The action of the light inflaton model augmented with three sterile neutrinos is \cite{Bezrukov:2009yw,Bezrukov:2013fca} \begin{align} S_{X\mathrm{SM}} = & \int\!\!\sqrt{-g}\,d^4x\left( \cL_\mathrm{SM} + \cL_{X H} + \cL_{N} + \cL_\text{grav} \right) , \nonumber \\ \cL_{XH} = & \frac{(\partial_\mu X)^2}{2} + \frac{m_X^2 X^2}{2} - \frac{\beta X^4}{4} - \lambda \left( H^\dagger H - \frac{\alpha}{\lambda} X^2 \right)^2 , \label{4*} \\ \cL_\text{grav} = & - \frac{M_P^2+\xi X^2}{2} R , \label{44} \\ \cL_{N} = & i\bar N_I\slashed\partial N_I - \left( F_{\alpha I}\bar L_\alpha N_I\tilde H +\frac{f_I}{2}\bar N_I^c N_I X+\mathrm{h.c.} \right) , \label{4+} \end{align} where $R$ is the scalar curvature, $\cL_\text{SM}$ is the SM Lagrangian without the EBH field potential, and $\cL_{N}$ stands for the renormalizable extension of the SM by 3 sterile neutrinos $N_I$ ($I=1,2,3$), $L_\alpha$ ($\alpha=e,\mu,\tau$) being lepton doublets and $\tilde H=\epsilon H^*$, where $\epsilon$ is $2\times2$ antisymmetric matrix and $H$ is the EBH doublet. With potential \eqref{4*} inflaton field $X$ gets vacuum expectation value, which breaks scale invariance both in the sterile neutrino sector (making sterile and active neutrino massive via \eqref{4+}) and in the SM sector (giving vacuum expectation $v$ to the EBH field via mixing in the last term of \eqref{4*}). Four parameters of the model, $m_X$, $\beta$, $\lambda$, and $\alpha$, determine the EBH field vacuum expectation value $v\approx246\GeV$, the Higgs boson mass $m_h\approx126\GeV$ \cite{Aad:2012tfa,Chatrchyan:2012ufa}, and the inflaton mass \begin{equation} \label{eq:10} m_\chi = m_h \sqrt{\frac{\beta}{2\alpha}} = \sqrt{\frac{\beta}{\lambda\theta^2}} . \end{equation} Thus, at a given value of $\beta$, the only free parameter in the scalar sector is the mixing coupling $\alpha$ or the inflaton mass $m_\chi$. The particle spectrum in vacuum consists of the Higgs boson $h$ and the inflaton $\chi$ of the mass $m_\chi$, which are mixed (as compared to the $H-X$ basis) by a small mixing angle \begin{equation} \label{eq:13} \theta^2 = \frac{2\beta v^2}{m_\chi^2} = \frac{2\alpha}{\lambda} . \end{equation} Hence, the branching ratios of the inflaton decay into the SM particles (see \cite{Bezrukov:2009yw} for details) are fixed for a given inflaton mass, see Fig.~\ref{fig:decays}. \begin{figure}[!htb] \centering \includegraphics{branchings} \caption{Inflaton decay branching rates \cite{Bezrukov:2009yw} for two-body final states. In the mass region $m_\chi\simeq 1\GeV$ predictions are highly uncertain because of the QCD effects.} \label{fig:decays} \end{figure} At large field values potential \eqref{4*} exhibits slow roll behavior along the direction $H^\dagger H=\alpha X^2/\lambda$, and supports the inflationary expansion of the early Universe. The non-minimal coupling to gravity \eqref{44} allows to control the amount of gravity waves generated at inflation \cite{Tsujikawa:2004my} and for $\xi\gtrsim10^{-3}$ is consistent \cite{Bezrukov:2013fca} with the Planck bounds \cite{Ade:2013uln}. The tilt of the scalar perturbation power spectrum also agrees with cosmological data. For a given $\xi$ the inflaton self-coupling $\beta$ is determined from the amplitude of the primordial density perturbations \cite{Bezrukov:2013fca}. To summarize, the model (\ref{4*},\ref{44}) has three new parameters $\xi$, $\beta$, and $m_\chi$ (or, equivalently, $\alpha$), in addition to SM \footnote{Parameter $m_X$ can be traded for the SM parameter $v$}, and they are determined from the following effects. \begin{itemize} \item $\beta$ and $\xi$ are related from the CMB normalization. \item $m_\chi$ and $\beta$ are related by the requirement of the generation of proper abundance of DM (given DM mass $M_1$ or coupling $f_1$ is known). \item $\xi$ can be determined from the measurement of the tensor-to-scalar ratio $r$ of the primordial perturbations. \end{itemize} This in principle completely fixes all the parameters of the model. We treat $m_\chi$ as free parameter in this \emph{Letter}, as far as the errors for the $r$ determination are still quite large, and will discuss this further in the Conclusions. At the same time the resulting values of the parameters should satisfy the set of constraints \cite{Bezrukov:2009yw,Bezrukov:2013fca} \begin{itemize} \item $\alpha$ is bound from below from the requirement of sufficient reheating, \item $\alpha$ is bound from above not to spoil the inflationary potential by radiative corrections, \item certain region in $m_\chi$ and $\theta$ (or, equivalently $\beta$) is constrained from particle physics experiments. \end{itemize} We show below that the first two are automatically satisfied with the parameters, leading to the proper DM generation, and the latter one leads to significant bound on the inflaton mass $m_\chi$ (and hence effective upper bound on $r$).
In the model, where the non-minimally coupled inflaton serves as the only scale invariance breaking messenger for SM with three sterile neutrinos, sterile neutrino DM can be generated in the inflaton decays\footnote{Two other sterile neutrinos giving masses to active neutrinos may be also adopted to explain baryon asymmetry of the Universe like in $\nu$MSM.}. In bosonic sector the model introduces three additional parameters---non-minimal coupling to gravity $\xi$, inflaton self coupling $\beta$, and the inflaton mass $m_\chi$. The amplitude of the primordial perturbations relates first two of these parameters, $\xi$ and $\beta$. Requirement of the proper abundance of the DM with a given mass $M_1$ provides the relation between the second pair, $\beta$ and $m_\chi$. Thus, assuming the mass of the DM is known, the only free parameter left is the inflaton mass. For numerical estimates we take $M_1=7\keV$, motivated by recent results \cite{Bulbul:2014sua,Boyarsky:2014jta}. Further constraints on the model can be made from inflationary observations, specifically the tensor-to-scalar ratio $r$. Exact knowledge of $r$ would fix the value of $\xi$, and, thus, the inflaton mass $m_\chi$ leaving no free parameters (see Figures~\ref{fig:particle-parameters2} and \ref{fig:cosmo-parameters}). Recent observation of $r\sim 0.2$ by BICEP2 \cite{Ade:2014xna} is in some tension with the CHARM bound on our model. However, if the dust foregrounds give major contribution to the BICEP2 signal, $r$ may be smaller and agree with predictions of our model. Thus, improvement or verification of the CHARM bound and independent checks of the BICEP2 result by multifrequency experiments are essential. Note also, that for heavier DM, $M_1>7\keV$, larger $r\sim0.2$ become allowed (cf.\ Figure~\ref{fig:particle-parameters2} and eq.\ (\ref{beta-for-mchi})). The resulting low mass range is especially interesting, as far as for the inflaton masses of $230\MeV\lesssim m_\chi\lesssim600\MeV$ the inflaton can be produced and searched in B-meson decays, see decay rate and inflaton lifetime in Figure~\ref{fig:particle-features}. For the lower masses the lifetime of the inflaton is relatively long and the most interesting signature is the offset vertex of the inflaton decay into muon or pion pair after the B-meson decay vertex. For the higher masses the inflaton lifetime drops rapidly, and the possible signature is the peak in the B-meson three body decay kinematics. Note that the expected event rates are comparable with the current experimental sensitivity \cite{Bezrukov:2013fca}. The lightest allowed inflatons may be searched for in beam-dump experiments (e.g. see SHIP {\bf REF} for CHARM successor). Finally, we should note, that there are ways to slightly relax the relations for the model parameters. First is the possibility of the entropy generation in the decays of the heavier sterile neutrinos after the DM generation ($S>1$). This allows for slightly larger $\beta$ for given $m_\chi$, leading to larger Higgs-inflaton mixing $\theta^2$, larger $\xi$, and smaller $r$. Additional generation of DM sterile neutrino $N_1$ after the inflaton decay (as in \cite{Boyarsky:2009ix}) leads to the opposite effect. More significant deviations is possible if we allow for additional sources of the violations of the scale invariance. Specifically, allowing for arbitrary mass terms for the sterile neutrinos independent of the inflaton coupling would allow to relax all the relations for the DM generation. If 7\keV{} sterile neutrino is not the dominant component of DM, $\Omega_N<\Omega_{DM}$, Higgs-inflaton mixing $\theta^2$ is smaller, hence larger $r$ is allowed. It is worth to study thoroughly the consistency of the particular mechanism of DM production with observations of Ly-$\alpha$ forest, since present lower limits on the DM free-streaming \cite{Boyarsky:2008xj,Viel:2013fqw} fall in the right ballpark. \bigskip The authors would like to thank M.~Shaposhnikov for valuable discussions. The work of D.G. is partly supported by RFBR grants 13-02-01127a and 14-02-00894a.
14
3
1403.4638
We study the mechanism of generation of 7 keV sterile neutrino Dark Matter (DM) in the model with light inflaton χ, which serves as a messenger of scale invariance breaking. In this model the inflaton, in addition to providing reheating to the Standard Model (SM) particles, decays directly into sterile neutrinos. The latter are responsible for the active neutrino oscillations via seesaw type I mechanism. While the two sterile neutrinos may also produce the lepton asymmetry in the primordial plasma and hence explain the baryon asymmetry of the Universe, the third one being the lightest may be of 7 keV and serve as DM. For this mechanism to work, the mass of the inflaton is bound to be light (0.1-1 GeV) and uniquely determines its properties, which allows to test the model. For particle physics experiments these are: inflaton lifetime (10<SUP>-5</SUP>-10<SUP>-12</SUP> s), branching ratio of B-meson to kaon and inflaton (10<SUP>-6</SUP>-10<SUP>-4</SUP>) and inflaton branching ratios into light SM particles like it would be for the SM Higgs boson of the same mass. For cosmological experiments these are: spectral index of scalar perturbations (n<SUB>s</SUB> ≃ 0.957- 0.967), and amount of tensor perturbations produced at inflation (tensor-to-scalar ratio r ≃ 0.15- 0.005).
false
[ "light inflaton χ", "light SM particles", "inflaton lifetime", "scale invariance breaking", "inflaton", "sterile neutrinos", "SM", "tensor perturbations", "ratio", "ratios", "particle physics experiments", "DM", "scalar perturbations", "Dark Matter", "keV", "7 keV sterile neutrino", "the SM Higgs boson", "third", "Universe", "kaon" ]
9.467936
-1.395202
77
437293
[ "Belluzzi, L.", "Trujillo Bueno, J." ]
2014A&A...564A..16B
[ "The transfer of resonance line polarization with partial frequency redistribution and J-state interference. Theoretical approach and numerical methods" ]
30
[ "Instituto de Astrofísica de Canarias, 38205, La Laguna, Tenerife, Spain ; Departamento de Astrofísica, Facultad de Física, Universidad de La Laguna, 38206, La Laguna, Tenerife, Spain", "Instituto de Astrofísica de Canarias, 38205, La Laguna, Tenerife, Spain; Departamento de Astrofísica, Facultad de Física, Universidad de La Laguna, 38206, La Laguna, Tenerife, Spain; Consejo Superior de Investigaciones Científicas, 28006, Madrid, Spain" ]
[ "2014ASPC..489..137T", "2015ApJ...812...73B", "2015ApJ...814..116B", "2015IAUS..305..127T", "2016A&A...591A..59B", "2016ApJ...824..135C", "2016ApJ...830L..24D", "2016ApJ...831L..15A", "2016IAUTA..29..278C", "2017A&A...607A..50B", "2017ApJ...836....6A", "2017MNRAS.470.1453Q", "2017SSRv..210..183T", "2018smf..book..183T", "2019A&A...624A.104J", "2019ApJ...880...85A", "2020ApJ...891...91D", "2021A&A...655A..13J", "2021A&A...655A..87J", "2021arXiv211011885J", "2022A&A...664A..76A", "2022A&A...664A.197B", "2022ARA&A..60..415T", "2022ApJ...933..145L", "2023ApJ...942...60A", "2023ApJ...954..218A", "2023ApJ...958...38J", "2024A&A...682A..68J", "2024A&A...683A.207G", "2024arXiv240117834R" ]
[ "astronomy" ]
8
[ "atomic processes", "line: formation", "polarization", "radiative transfer", "scattering", "stars: atmospheres", "Astrophysics - Solar and Stellar Astrophysics" ]
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[ "10.1051/0004-6361/201321598", "10.48550/arXiv.1403.1701" ]
1403
1403.1701_arXiv.txt
To probe the thermal, dynamic and magnetic properties of the chromosphere and transition region of the Sun and of other stars, we need to measure and interpret the intensity and polarization profiles of strong resonance lines, such as Mg~{\sc ii} h \& k, hydrogen Ly-$\alpha$, or He~{\sc ii} 304~{\AA}. Interpretation of the Stokes parameters of the observed radiation requires solving a non-LTE radiative transfer problem that can be very complex, especially when the main interest lies in modeling the spectral details of the linear polarization signals produced by scattering processes and their modification by the Hanle effect. One of the main difficulties occurs because scattering is a second-order process within the framework of quantum electrodynamics \citep[e.g.,][]{Cas07}, where frequency correlations between the incoming and outgoing photons can occur (partial frequency redistribution, PRD). Another aspect contributing to the complexity of the problem is the need to account for the impact of quantum interference (or coherence) between pairs of magnetic sublevels pertaining either to the same $J$-level ($J$ being the level's total angular momentum) or to different $J$-levels of the same term \citep[$J$-state interference; see][and references therein]{Bel11}. An additional complication stems from the fact that the plasma of a stellar chromosphere can be highly inhomogeneous and dynamic, which implies the need to solve the non-equilibrium problem of the generation and transfer of polarized radiation in realistic three-dimensional stellar atmospheric models \citep[e.g.,][]{Step13}. As shown by \citet{Bel12a}, the joint action of PRD and $J$-state interference produces a complex scattering polarization $Q/I$ profile across the h and k lines of Mg~{\sc ii}, with large amplitudes in their wings. The same happens with the Ly-$\alpha$ lines of hydrogen and ionized helium \citep{Bel12b}. In these {\it Letters to the Editor} we only had room to briefly outline our approach to the problem and to present the main results. The aim of this paper is therefore to describe the theoretical framework and the numerical methods we have developed for modeling the transfer of resonance line polarization taking PRD and $J$-state interference into account. An atomic model accounting for the presence of quantum interference between pairs of sublevels pertaining either to the same $J$-level or to different $J$-levels within the same term is the {\it multiterm} model atom described in the monograph by Landi Degl'Innocenti \& Landolfi (2004; hereafter LL04). While a quantum-mechanical PRD theory for a two-level atom is available \citep[see][]{Bom97a,Bom97b}, a self-consistent PRD theory for multiterm systems has not been derived yet. The approximate approach that we propose in this paper is heuristically built starting from two distinct theoretical frameworks: the polarization theory presented in LL04 and the theoretical scheme described in \citet{Lan97}. The density matrix theory of LL04 is based on a lowest order perturbative expansion of the atom-photon interaction within the framework of quantum electrodynamics. Within this theoretical approach, a scattering event is described as a temporal succession of independent first-order absorption and re-emission processes, which is justified when the plasma conditions are such that the coherency of scattering is completely destroyed (limit of complete redistribution in frequency, CRD). The theoretical approach developed by \citet{Lan97} allows us to describe the opposite limit of purely coherent scattering in the atom rest frame. This theory is based on the assumption that the atomic energy levels are composed by a continuous distribution of so-called metalevels \citep[see,][]{Hub83a,Hub83b}. In \citet{Lan97}, the expression of the redistribution matrix for coherent scattering \citep[$R_{II}$, following the terminology introduced by][]{Hum62} is derived for different atomic models. In particular, the authors consider the case of a two-term atom, under the assumptions that the fine structure levels of the lower term are infinitely sharp, and that the magnetic sublevels of the lower term are evenly populated and no interference is present between them (unpolarized lower term). We consider a two-term model atom with unpolarized lower term and infinitely sharp lower levels. We point out that this atomic model is not just of academic interest, but it is perfectly suitable for investigating several resonance doublets of high diagnostic interest, such as the h \& k lines of Mg~{\sc ii}, or the Ly-$\alpha$ lines of hydrogen and ionized helium. This is because the lower level of these lines has $J=1/2$ (it cannot be polarized, unless the incident radiation has net circular polarization) and it is the ground level (its long lifetime justifies the infinitely-sharp level assumption). By analogy with the two-level atom case (briefly recalled in Sect.~2), we assume that in the atom rest frame the total redistribution matrix for such a two-term atom is given by a linear combination of two terms, one describing the limit of coherent scattering ($R_{II}$) and one describing the CRD limit ($R_{III}$). Starting from the two-term atom theory presented in LL04, as generalized by \citet{Bel13a} for including the effect of inelastic and superelastic collisions\footnote{Hereafter, following LL04, we will distinguish between inelastic and superelastic collisions. The former induce transitions towards higher energy levels (`exciting' collisions), the latter induce transitions towards lower energy levels (`de-exciting' collisions).}, we derive an approximate expression of the $R_{III}$ redistribution matrix, under the assumption of CRD in the observer's frame (Sect.~3.1). Concerning the $R_{II}$ redistribution matrix, we start from the expression derived by \citet{Lan97} in the atomic rest frame. We derive the corresponding expressions in the observer's frame taking Doppler redistribution into account, and we include the effect of inelastic and superelastic collisions by analogy with the case of $R_{III}$ (Sect.~3.2). The formulation of the radiative transfer problem is presented in Sect.~5, while the iterative method for the solution of the non-LTE problem is discussed in Sect.~6. An illustrative application of our modeling scheme is shown in Sect.~7.
\label{Sect:conclusions} Information on the magnetism of the outer atmosphere of the Sun and of other stars is encoded in the polarization that several physical mechanisms introduce in spectral lines. Of particular interest is the modification, due to the Hanle effect induced by weak ($B{\lesssim}$100~G) magnetic fields, of the linear polarization signals produced by scattering processes in strong resonance lines. In order to model correctly the spectral details of scattering polarization signals generated in strong resonance lines in an optically thick plasma, it is essential to solve the ensuing non-LTE radiative transfer problem taking PRD and $J$-state interference effects into account. In this paper we have presented a theoretical approach to this complex problem, as well as the numerical methods for the solution of the ensuing equations. This first investigation has been restricted to the unmagnetized reference case. We have considered a two-term atom with unpolarized lower term and infinitely sharp lower levels. This atomic model accounts for quantum interference between sublevels pertaining either to the same $J$-level or to different $J$-levels of the upper term, and it is suitable for modeling several resonance lines of high diagnostic interest (e.g., Mg~{\sc ii} h and k, H~{\sc i} Ly-$\alpha$, He~{\sc ii} Ly-$\alpha$). By analogy with the two-level atom case, we have assumed that the redistribution matrix for such an atomic model is given by the linear combination of two terms, one describing coherent scattering processes in the atom rest frame ($R_{II}$), and one describing scattering processes in the limit of complete frequency redistribution ($R_{III}$). The expression of $R_{II}$ has been derived within the framework of the theoretical approach presented in \citet{Lan97} (suitable for treating the limit of coherent scattering), while the expression of $R_{III}$ has been deduced within the framework of the theory of polarization described in LL04 (suitable for treating the limit of CRD). Concerning the $R_{III}$ redistribution matrix, we have derived an approximate expression that describes scattering events in the limit in which collisions are able to redistribute the photon frequency across the whole multiplet. As discussed in Sect.~3.1, this is in general a strong assumption, especially when the energy separation among the various $J$-levels is very large. This assumption is ultimately due to the fact that our $R_{III}$ has been derived in a heuristic way starting from the theory of LL04, which is based on the flat-spectrum approximation. Although other approximate forms of $R_{III}$, based on different hypotheses, can in principle be proposed \citep[see, for example,][]{Smi13}, we believe that the theory of LL04 is at the moment the most robust one for deriving this redistribution matrix. A comparison between the different expressions of $R_{III}$ that have been proposed up to now, and a detailed analysis of their properties and weaknesses will be the object of a forthcoming paper. The depolarizing effect of elastic and weakly inelastic collisions has been neglected. This is not fully consistent, but accounting for this effect gives rise to a series of difficulties that go beyond the scope of this work. As discussed in Sect.~3.3, this should be in any case a suitable approximation for modeling the line-core signal of strong resonance lines forming in the upper chromosphere and transition region of the Sun. As pointed out in Sect.~4, performing a series of formal substitutions of the quantum numbers, the redistribution matrices presented in this work for a two-term atom can be applied to describe a two-level atom with HFS (i.e., a model atom accounting for interference between pairs of magnetic sublevels pertaining either to the same HFS $F$-level, or to different $F$-levels of the same $J$-level). The hypotheses under which the physical scenario described by our $R_{III}$ redistribution matrix is strictly justified turn out to be less restrictive in a two-level atom with HFS than in a two-term atom. The numerical method of solution we have developed for solving this non-LTE radiative transfer problem is based on a direct generalization of the Jacobian iterative scheme developed by \citet{JTB99} for the case of a two-level atom (i.e., without $J$-state interference), in the limit of CRD. The ensuing computer program can be applied to investigate a number of interesting radiative transfer problems in solar and stellar physics \citep[e.g.,][]{Bel12a,Bel12b,Bel13b}. Moreover, it can be used as starting point for further developments, such as the modeling of the Hanle effect and/or the investigation of the scattering polarization that results from the interaction between the line and continuum processes. We plan to address several of these issues in forthcoming papers.
14
3
1403.1701
The linear polarization signals produced by scattering processes in strong resonance lines are rich in information on the magnetic and thermal structure of the chromosphere and transition region of the Sun and of other stars. A correct modeling of these signals requires accounting for partial frequency redistribution effects, as well as for the impact of quantum interference between different fine structure levels (J-state interference). In this paper, we present a theoretical approach suitable for modeling the transfer of resonance line polarization when taking these effects into account, along with an accurate numerical method of solution of the problem's equations. We consider a two-term atom with unpolarized lower term and infinitely sharp lower levels, in the absence of magnetic fields. We show that by making simple formal substitutions on the quantum numbers, the theoretical approach derived here for a two-term atom can also be applied to describe a two-level atom with hyperfine structure. An illustrative application to the Mg ii doublet around 2800 Å is presented.
false
[ "different fine structure levels", "hyperfine structure", "other stars", "quantum interference", "resonance line polarization", "strong resonance lines", "magnetic fields", "unpolarized lower term", "partial frequency redistribution effects", "J-state interference", "Sun", "infinitely sharp lower levels", "simple formal substitutions", "the magnetic and thermal structure", "solution", "information", "account", "an accurate numerical method", "the chromosphere and transition region", "a theoretical approach" ]
11.936462
14.729097
2
483090
[ "Schönrich, Ralph", "Asplund, Martin", "Casagrande, Luca" ]
2014ApJ...786....7S
[ "Does SEGUE/SDSS indicate a dual Galactic halo?" ]
35
[ "Department of Astronomy, The Ohio State University, 140 West 18th Avenue, Columbus, OH 43210, USA ; Hubble Fellow.;", "Research School of Astronomy and Astrophysics, Australian National University, Cotter Road, Weston ACT 2611, Canberra, Australia", "Research School of Astronomy and Astrophysics, Australian National University, Cotter Road, Weston ACT 2611, Canberra, Australia" ]
[ "2014A&A...568A...7A", "2014A&A...569A..13R", "2014ApJ...787...30D", "2014EAS....65..349H", "2014MNRAS.445L..21T", "2015A&A...577A..81F", "2015ARA&A..53..631F", "2015ApJ...813...89K", "2015MNRAS.446.2274L", "2015MNRAS.449.2604D", "2015MNRAS.452..676P", "2015MNRAS.452.2675B", "2016AN....337..889C", "2016ARA&A..54..529B", "2016ASSL..435...93R", "2016MNRAS.455..987C", "2016MNRAS.460.1725D", "2016RAA....16...44G", "2017ASSL..434...31C", "2017ApJ...845..101B", "2017MNRAS.465..501S", "2017MNRAS.470.2959K", "2018A&A...615A..70P", "2018ApJ...852..128G", "2019ApJ...882..176K", "2019ApJ...885..102L", "2019IAUS..339...50E", "2019MNRAS.484..218G", "2019MNRAS.487.1462F", "2019RAA....19....8G", "2020AJ....159..270K", "2020MNRAS.492.3631M", "2020MNRAS.492.4986Y", "2020MNRAS.493.5195D", "2023MNRAS.520.1832B" ]
[ "astronomy" ]
6
[ "galaxies: halos", "Galaxy: halo", "Galaxy: kinematics and dynamics", "Galaxy: structure", "solar neighborhood", "stars: distances", "Astrophysics - Astrophysics of Galaxies" ]
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[ "10.1088/0004-637X/786/1/7", "10.48550/arXiv.1403.0937" ]
1403
1403.0937_arXiv.txt
Since the discovery of the Galactic halo \citep[][]{Schwarzschild52, Eggen62}, its structure and origin have been under intense debate. Three main sources of halo stars have been discussed in the past \citep[see][ for a discussion]{Sheffield12}: monolithic collapse in early galaxy formation as put forth by \cite{Eggen62}; the accretion of satellites \citep[][]{Searle78}; and kick-up of disc/bulge stars via minor mergers \citep[e.g.][]{Purcell10}, ejections from cluster cores \citep[e.g.][]{Leonard91}, or binary interactions \citep[][]{Przybilla08} involving a supermassive or intermediate mass black hole \citep[e.g.][]{Hills88, Gualandris07}. Of these the second is considered to be the dominant contribution to the Galactic halo based on $\Lambda CDM$ cosmological simulations and from observational data \citep[see][]{Sheffield12}, The impact of kicked out stars is strongly limited theoretically, as well as observationally by the low number of metal-rich halo stars. It can be hypothesized that stars from an initial collapse should on average carry some prograde momentum. For accreted stars, early studies \citep[][]{Quinn86, Byrd86}, recently confirmed by \citep[][]{Murante10}, suggested that dynamic friction differentiates between prograde and retrograde infalling satellites, leaving behind a potentially detectable asymmetry in kinematics. It is, however, not clear, how this signature should translate into different rotation as a function of metallicity. \cite[][]{Cooper10} found that the accretion signatures along with radial abundance gradients vary strongly in simulated galaxies, depending strongly on the individual accretion history. In principle three different observational signatures have been proposed: \begin{itemize} \item Asymmetry in the (azimuthal) velocity distribution, which may be a consequence of dynamic friction. However, other processes as well as the special accretion history can lead to such a distribution. This dates back to \cite[][]{Norris89}, but it has been known since \cite{Strom27} and later \cite{Ryan92} that distance errors can account for this type of asymmetries. \item Radial metallicity gradients. As noted before, these depend on the specific accretion history of a Galaxy. By tendency the material accreted later onto a galaxy stems from less dense regions that evolved later and produced less massive stellar systems \citep[this was already pointed out by][]{Kant1755} and hence are likely to have a lower metallicity. Claims of related differences between the inner and outer halo were made as early as \cite{Searle78} and \cite{Preston90}. \item Radial gradients in angular momentum/ mean azimuthal velocity. Again, this hinges on the detailed accretion history, correlating velocities and metallicities to varying degree and sign. For our Galaxy, BHB stars set quite narrow limits on any radial trend \citep[][]{FSII}. \end{itemize} \cite{Carollo07} and \cite{Carollo10} (hereafter C07 and C10) claimed the existence of a dual halo that consists of a prograde, metal-rich ($\feh \approx -1.6$), inner halo and a distinct retrograde, more metal-poor ($\feh \approx -2.2$), outer halo. This observational result triggered a major series of theoretical papers finding such structures in simulations \citep[e.g.][]{Zolotov09, Zolotov10, Font11, McCarthy12, Tissera12}, while some other studies \citep[e.g.][]{Lucia08, Cooper10} find no convincing trends. The result was criticized by \cite{SAC} (hereafter SAC11),tracing their claims back to the inappropriate use of Gaussian analysis and the neglect of observational errors, as well as unphysical distance estimates in their sample. Accounting for these effects, SAC11 showed that on any reasonable adopted distance scale, including C10's own distances, their findings of a halo duality vanished. Recently \cite{Beers12} (hereafter B12) published a re-analysis affirmative of the original C07,10 studies, finding again a dual halo with a retrograde metal-poor component. It could be argued that this question will soon be answered by the Gaia satellite mission. While this is likely true, the purpose of the present paper is more general, i.e. to lay out on the example of B12 the many subtleties due to selection biases and error analysis, which all studies of this kind even with Gaia data need to address in order to derive meaningful conclusions. With this spirit we make the data used in this work publicly available\footnote{Please find the data under \newline {\it \scriptsize http://www-thphys.physics.ox.ac.uk/people/RalphSchoenrich/data/halo/}. We will be delighted to provide any additional information upon request.} to enable independent investigations on the topic by others.
\label{sec:conclude} We have studied the arguments presented by B12 for a dual halo with a retrograde, metal-poor outer component and conclude that none of them stands up to critical examination. In the course of our replication of the B12 analysis we found that the relations from \cite{Beers00} are significantly brighter than stellar models in the Johnson V-band. In addition we checked their performance on M92 and could not reconcile the Beers et al. distances with the available data. Hence we caution against use of the calibrations of \cite{Beers00} for metal-poor stars. We also point out that B12 wrongly ascribed to us and named after us a distance relation that is far shorter/fainter than anything we ever used, just to show that it is too faint. The trend of metallicity with altitude claimed by B12 is primarily driven by disc stars. The smaller trend for low metallicities that exclude those disc objects is identified as a selection bias: metal-poor stars are more luminous and in SEGUE are more likely identified as brighter subgiant and turn-off stars in this sample. Hence the metallicity has a spurious correlation with altitude via its dependence on distance. Controlling for this effect, the altitude dependence of halo star metallicities vanishes. We identified an analogous problem for BHB stars and found a likely contamination with disc objects. The selection bias investigated in Section \ref{sec:ddd} is an excellent example of why forward modelling from theoretical models is needed and why measured distributions from observations in a non-volume complete sample are unreliable without a thorough analysis and correction of biases. The difference of proper motion distributions between retrograde stars and the remaining sample is not an argument that their data must be reliable, since stars with zero proper motions require extremely high line-of-sight velocities to be retrograde (and a position near the Galactic azimuth): retrograde stars essentially must show large proper motions. In fact, the contrast between ''retrograde`` (as measured) stars and the remaining population sharpens in proper motion space with decreasing sample quality. Furthermore, B12 used unphysical Gaussian analysis on the azimuthal velocity distribution. As emphasized in SAC11, such analysis is inappropriate as the observational error is highly non-Gaussian and has long tails, especially towards the retrograde side. We demonstrated - now on the highly metal-poor stars - that we can easily fit it by an underlying Gaussian distribution folded with the appropriate measurement errors. The B12 finding of a difference between the velocity distributions of BHB stars does not trace back to a different behaviour between metallicities as they argue, but to hot metal-poor BHB stars behaving differently from their less hot metal-poor counterparts. Further checks do not reveal a reliable retrograde signal. Further evidence in the literature is either subject to similar statistical and selection biases or not specific to the claim of a dual halo. For example, neither does the presence of increased substructure in the outskirts of the Galactic halo (which is a natural consequence of accretion and dynamic friction) constitute an indication for a retrograde, metal-poor outer halo, nor does the theoretical possibility of a duality as suggested by some (but far from all) simulations act as any observational evidence. While we find no evidence for any chemokinematic halo duality in the SEGUE sample, this disproof of evidence should not be misinterpreted as a proof of absence. Even with bold extrapolation the sample does not allow investigations beyond a distance of roughly $15 \kpc$. Further data and investigations especially with more remote samples, like the BOSS \citep[Baryon Oscillation Spectroscopic Survey][]{Dawson13} spectra, should be gathered to search for more subtle trends or trends beyond the current distance limits. We repeat our statement from SAC11 that we favour a halo that shows an increase of substructure towards larger Galactocentric radii and do not try to address the question if there is any metallicity gradient towards its outskirts beyond what is probed by SEGUE. What we can definitely say is that there is no evidence from current SEGUE data that would support the claims of C07, C10 or B12 of a dual halo with a retrograde, highly metal-poor outer halo.
14
3
1403.0937
We re-examine recent claims of observational evidence for a dual Galactic halo in SEGUE/SDSS data, and trace them back to improper error treatment and neglect of selection effects. In particular, the detection of a vertical abundance gradient in the halo can be explained as a metallicity bias in distance. A similar bias and the impact of disk contamination affect the sample of blue horizontal branch stars. These examples highlight why non-volume complete samples require forward modeling from theoretical models or extensive bias-corrections. We also show how observational uncertainties produce the specific non-Gaussianity in the observed azimuthal velocity distribution of halo stars, which can be erroneously identified as two Gaussian components. A single kinematic component yields an excellent fit to the observed data, when we model the measurement process including distance uncertainties. Furthermore, we show that sample differences in proper motion space are the direct consequence of kinematic cuts and are enhanced when distance estimates are less accurate. Thus, their presence is neither proof of a separate population nor a measure of reliability for the applied distances. We conclude that currently there is no evidence from SEGUE/SDSS that would favor a dual Galactic halo over a single halo that is full of substructure.
false
[ "distance uncertainties", "distance estimates", "halo stars", "distance", "non-volume complete samples", "blue horizontal branch stars", "selection effects", "observational uncertainties", "improper error treatment", "sample differences", "theoretical models", "Gaussian", "disk contamination", "kinematic cuts", "proper motion space", "a dual Galactic halo", "A single kinematic component", "observational evidence", "two Gaussian components", "the applied distances" ]
9.389338
7.466416
174
814229
[ "Kofinas, Georgios", "Saridakis, Emmanuel N.", "Xia, Jun-Qing" ]
2014PhRvD..90h4049K
[ "Cosmological solutions and observational constraints on five-dimensional braneworld cosmology with gravitating Nambu-Goto matching conditions" ]
2
[ "Department of Information and Communication Systems Engineering, Research Group of Geometry, Dynamical Systems and Cosmology, University of the Aegean, Karlovassi 83200, Samos, Greece", "Physics Division, National Technical University of Athens, 15780 Zografou Campus, Athens, Greece; Instituto de Física, Pontificia Universidad Católica de Valparaíso, Casilla 4950, Valparaíso, Chile", "Key Laboratory of Particle Astrophysics, Institute of High Energy Physics, Chinese Academy of Science, P.O. Box 918-3, Beijing 100049, China" ]
[ "2015IJMPD..2430014H", "2018PhLB..780...29D" ]
[ "astronomy", "physics" ]
3
[ "04.50.-h", "95.36.+x", "98.80.-k", "Higher-dimensional gravity and other theories of gravity", "Dark energy", "Cosmology", "Astrophysics - Cosmology and Nongalactic Astrophysics", "General Relativity and Quantum Cosmology", "High Energy Physics - Phenomenology", "High Energy Physics - Theory" ]
[ "1966NCimB..44....1I", "1967NCimB..48..463I", "1971JMP....12..498L", "1975JMP....16...65G", "1977PhRvD..15..935I", "1981PhRvD..23..852V", "1985PhR...121..263V", "1986PhR...137..109Z", "1987CQGra...4....1V", "1987PhRvD..36..392M", "1987PhRvD..36.1017G", "1989PhRvD..39.1084F", "1989PhRvL..62.2897U", "1990CQGra...7....1C", "1993PhR...229..145M", "1999CQGra..16.4101G", "1999PhRvL..83.3370R", "1999PhRvL..83.4690R", "2000PhLB..477..285B", "2002PhRvD..66d3521I", "2002PhRvL..88w1101G", "2003JHEP...06..048C", "2003PThPh.110..201N", "2003PhLB..562..118G", "2003PhRvD..67b4030D", "2004AnPhy.309..232E", "2004PhRvL..92v1601B", "2005ApJ...633..560E", "2005JHEP...08..075C", "2005PhRvD..72d4013A", "2006PhRvD..74d4023D", "2007CQGra..24.5861K", "2008JCAP...07..012L", "2010JCAP...01..013D", "2010JCAP...01..022C", "2012ApJ...746...85S", "2014A&A...571A..16P", "2014AnPhy.351..504K", "2014PhRvD..89f5015K" ]
[ "10.1103/PhysRevD.90.084049", "10.48550/arXiv.1403.7510" ]
1403
1403.7510_arXiv.txt
\label{Introduction} Israel matching conditions \cite{Israel 1966} are considered as the standard equations of motion of a classical codimension-1 defect which backreacts on the bulk dynamics. They are derived by focusing on the distributional part of the Einstein field equations (or some gravity modification) where the brane energy-momentum tensor, specified by a delta function, is included. An equivalent way to derive these equations is to take the variation of the brane-bulk action with respect to the induced metric, while the bulk equations of motion are derived as usually by varying the bulk action with respect to the bulk metric. However, a higher codimension defect carrying a generic energy-momentum tensor is known to be inconsistent with Einstein's equations \cite{Israel 1977,Geroch,Garfinkle:1999xv} (a brane with a pure tension is a special consistent case \cite{Vilenkin:1981zs,VilenkinPR,Vickers:1987az,Frolov:1989er,Unruh:1989hy, Clarke:1990hu, Nakamura:2003pi,Cline:2003ak,Apostolopoulos:2005ff}). In \cite{Bostock:2003cv} it was considered the idea that a more general theory like Einstein-Gauss-Bonnet gravity in six dimensions might remove the previous inconsistency, and the matching conditions of the theory for a generic energy-momentum tensor were derived. In \cite{CKP} the consistency of the whole system of bulk field equations plus matching conditions was shown for an axially symmetric codimension-2 cosmological brane. The spirit of the above proposal for consistency of the higher codimension defects is to include higher Lovelock densities \cite{Lovelock:1971yv,Zumino:1985dp}. However, in $D$ dimensions, the highest such density is of order $[(D-1)/2]$, and so, it is quite probable that branes with codimension higher than $[(D-1)/2]$ will still be inconsistent. Moreover, four-dimensions which represent effectively spacetime at certain length and energy scales do not allow generic codimension-2 or 3 defects. On the other hand, Israel matching conditions and their generalizations to higher codimensions do not accept the Nambu-Goto probe limit, which is the lowest order approximation of a test brane moving in curved background. Even the geodesic limit of the Israel matching conditions is questionable as a probe limit, since being the geodesic equation a kinematical fact it should be preserved independent of the gravitational theory (similarly to \cite{GerochJang}, \cite{Ehlers:2003tv}) or the codimension of the defect, which is not the case for these matching conditions \cite{Bostock:2003cv,Germani,Davis:2002gn,Gravanis:2002wy,Myers,Charmousis}. Moreover, even the non-geodesic probe limit of the standard equations of motion for various codimension defects in Lovelock gravity theories is not accepted, since this consists of higher order algebraic equations in the extrinsic curvature, and therefore, a multiplicity of probe solutions arise instead of a unique equation of motion at the probe level. In view of these observations a criticism to the standard matching conditions appeared in \cite{kof-ira}, where alternative matching conditions were proposed. These are the ``gravitating Nambu-Goto matching conditions'' which arise by the variation of the brane-bulk action with respect to the brane embedding fields, so that the gravitational back-reaction of the brane is taken into account. With these matching conditions a brane is always consistent for an arbitrary energy-momentum tensor and it also possesses the Nambu-Goto probe limit (the codimension-2 case was studied in \cite{kof-ira}, \cite{KofTom}, while the codimension-1 in \cite{kof-zar}). In \cite{kof-zar} the application of these alternative matching conditions led to a new 5-dimensional braneworld cosmology which generalizes the conventional braneworld cosmology \cite{binetruy} in the sense that it contains an extra integration constant, and vanishing this constant gives back the standard braneworld cosmology. In the current work we try to confront this cosmology with the current cosmological observational data (SNIa, BAO, BBN) in order to construct the corresponding probability contour-plots for the parameters of the theory. The paper is organized as follows: In section \ref{Introduction of the model} we briefly present the alternative matching conditions and the basic features behind these, and we find in the cosmological framework the equation for the expansion rate including both the matter and radiation sectors. In section \ref{Observational}, which is the main part of the work, we impose the observational constraints on the parameters of the model. Finally, a summary of the obtained results is given in section \ref{Conclusions} of conclusions.
\label{Conclusions} In this work we constrained an alternative 5-dimensional braneworld cosmology using observational data. The difference with the standard braneworld cosmology refers to the adaptation of alternative matching conditions introduced in \cite{kof-ira} which generalize the conventional matching conditions. The reasons for this consideration are possible theoretical deficiencies of the standard junction conditions, namely the need for consistency of the various codimension defects and the existence of a meaningful equation of motion at the probe limit. Instead of varying the brane-bulk action with respect to the bulk metric at the brane position and derive the standard matching conditions, we vary with respect to the brane embedding fields in a way that takes into account the gravitational back-reaction of the brane onto the bulk. The proposed gravitating Nambu-Goto matching conditions may be close to the correct direction of finding realistic matching conditions since they always have the Nambu-Goto probe limit (independently of the gravity theory, the dimensionality of spacetime or codimensionality of the brane), and moreover, with these matching conditions, defects of any codimension seem to be consistent for any (second order) gravity theory. Compared to the conventional 5-dimensional braneworld cosmology, the new one possesses an extra integration constant, which if set to zero reduces the new cosmology to the conventional braneworld one. In the present work we extended the codimension-1 cosmology of \cite{kof-zar} by allowing both a matter and a radiation sector in order to extract observational constraints on the involved model parameters. In particular, we used data from supernovae type Ia (SNIa) and Baryon Acoustic Oscillations (BAO), along with arguments from Big Bang Nucleosynthesis (BBN) in order to construct the corresponding probability contour-plots for the parameters of the theory. Concerning the first ($\epsilon=-1$) branch of cosmology, we found that the parameters $\tilde{V}$ and $\tilde{c}$ that quantify the deviation from the Randall-Sundrum scenario, are constrained very close to their RS values as expected. However, a departure from Randall-Sundrum scenario is still allowed, and moreover, the corresponding $\chi^2$ is the same for both models. This means that braneworld models with gravitating Nambu-Goto matching condition are in ``equal'' agreement with observations with standard braneworld cosmology. However, application of the AIC criterion shows that both standard braneworld cosmology, as well as the extended scenario of the present work, are less favored by the data if we compare them with the concordance $\Lambda$CDM cosmology since the two extra parameters do not improve the fitting behavior. Furthermore, the obtained age of the universe is $12.23\ \text{Gyr}\leq t_0\leq14.13\ \text{Gyr}$, which is an additional observational advantage of the model. Finally, concerning the fundamental mass scale $M$, the current age estimations imply that the preferred values of $M$ lie well below the GeV scale. Concerning the second ($\epsilon=+1$) cosmological branch, which is completely new and with no correspondence in Randall-Sundrum scenario, we extracted the corresponding likelihood contours. Although this case is still compatible with observations, the corresponding minimal $\chi^2$ is much higher than that for the $\epsilon=-1$ branch case, which means that this branch case is not favored by late-times observations. However, although this branch is not favored by late-times observations, due to that $H^{2}\approx \text{const.}$ at early times, it could still play an important role in the inflationary regime. In summary, cosmology with gravitating Nambu-Goto matching conditions offers an extension to the standard Randall-Sundrum scenario. Apart from interesting solutions, we see that it is in agreement with observations since the data allow for a small deviation from Randall-Sundrum cosmology. Therefore, it should be worthy to further study the cosmological implications of the model, such as the inflationary behavior and the late-times asymptotic features, since especially a successful inflationary regime is something that cannot be obtained in the framework of $\Lambda$CDM cosmology.
14
3
1403.7510
We investigate the cosmological implications of the recently constructed five-dimensional braneworld cosmology with gravitating Nambu-Goto matching conditions. Inserting both matter and radiation sectors, we extract analytical cosmological solutions. Additionally, we use observational data from Type Ia supernovae and baryon acoustic oscillations, along with requirements of big bang nucleosynthesis and cosmic microwave background radiation, in order to impose constraints on the parameters of the model. We find that the scenario at hand is in good agreement with observations, and thus a small departure from the standard Randall-Sundrum scenario is allowed. However, the concordance Λ CDM cosmology is still favored comparing to both the standard braneworld model and the present scenario.
false
[ "good agreement", "cosmic microwave background radiation", "big bang nucleosynthesis", "analytical cosmological solutions", "observations", "the present scenario", "the standard Randall-Sundrum scenario", "Λ CDM cosmology", "hand", "Nambu-Goto matching conditions", "order", "constraints", "the scenario", "Type Ia supernovae and baryon acoustic oscillations", "both the standard braneworld model", "requirements", "Randall-Sundrum", "observational data", "the model", "thus a small departure" ]
10.182183
0.079822
-1
483563
[ "Tadaki, Ken-ichi", "Kodama, Tadayuki", "Tamura, Yoichi", "Hayashi, Masao", "Koyama, Yusei", "Shimakawa, Rhythm", "Tanaka, Ichi", "Kohno, Kotaro", "Hatsukade, Bunyo", "Suzuki, Kenta" ]
2014ApJ...788L..23T
[ "Evidence for a Gas-rich Major Merger in a Proto-cluster at z = 2.5" ]
25
[ "National Astronomical Observatory of Japan, Mitaka, Tokyo 181-8588, Japan", "National Astronomical Observatory of Japan, Mitaka, Tokyo 181-8588, Japan; Department of Astronomical Science, The Graduate University for Advanced Studies, Mitaka, Tokyo 181-8588, Japan", "Institute of Astronomy, The University of Tokyo, Mitaka, Tokyo 181-0015, Japan", "National Astronomical Observatory of Japan, Mitaka, Tokyo 181-8588, Japan; Institute for Cosmic Ray Research, The University of Tokyo, 5-1-5 Kashiwanoha, Kashiwa, Chiba 277-8582, Japan", "National Astronomical Observatory of Japan, Mitaka, Tokyo 181-8588, Japan; Institute of Space Astronomical Science, Japan Aerospace Exploration Agency, Sagamihara, Kanagawa 252-5210, Japan", "Department of Astronomical Science, The Graduate University for Advanced Studies, Mitaka, Tokyo 181-8588, Japan", "Subaru Telescope, National Astronomical Observatory of Japan, 650 North A'ohoku Place, Hilo, HI 96720, USA", "Institute of Astronomy, The University of Tokyo, Mitaka, Tokyo 181-0015, Japan; Research Center for the Early Universe, The University of Tokyo, 7-3-1 Hongo, Bunkyo, Tokyo 113-0033, Japan", "National Astronomical Observatory of Japan, Mitaka, Tokyo 181-8588, Japan", "Institute of Astronomy, The University of Tokyo, Mitaka, Tokyo 181-0015, Japan" ]
[ "2014ApJ...795....5O", "2014SPIE.9147E..96E", "2015ASPC..499...65K", "2015ApJ...800...67L", "2015MNRAS.448..666S", "2015MNRAS.449L..68C", "2016ApJ...830...89M", "2017A&A...608A..48D", "2017ApJ...841L..21H", "2017ApJ...842...55L", "2017MNRAS.468.1123S", "2017MNRAS.469.2235K", "2018ApJ...856..118H", "2018ApJ...862...64S", "2018ApJ...867L..29W", "2019A&A...625A..96K", "2019ApJ...883...92L", "2019PASJ...71...40T", "2021A&A...652A..11J", "2021arXiv211100813A", "2022A&A...662A..85P", "2022ApJ...924...74A", "2023MNRAS.523.2422L", "2023MNRAS.523.2738M", "2024arXiv240306729S" ]
[ "astronomy" ]
6
[ "galaxies: evolution", "galaxies: high-redshift", "galaxies: ISM", "Astrophysics - Astrophysics of Galaxies", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
[ "1955ApJ...121..161S", "1996ApJ...464..641M", "1997ApJ...478..144S", "1997ApJ...490..577D", "1998ARA&A..36..189K", "1998ApJ...506L...7F", "1998ApJ...507..615D", "1999MNRAS.302..152K", "2001ApJ...548..681B", "2002MNRAS.337....1I", "2003ApJ...597L.113N", "2005ASPC..347...81M", "2005MNRAS.359.1165G", "2007ASPC..376..127M", "2007MNRAS.377.1717K", "2007PASJ...59..117K", "2008AJ....136.2563W", "2008ApJ...680..246T", "2008ApJ...682..985W", "2008ApJ...687...59G", "2009Natur.457..451D", "2009PASJ...61..481S", "2010ApJ...713..686D", "2010ApJ...723.1139H", "2010ApJ...724..233E", "2010MNRAS.403.1611K", "2010MNRAS.407.2091G", "2011ApJ...730L..19G", "2011ApJ...737...67M", "2011ApJ...738..106W", "2011ApJ...740L..15M", "2011MNRAS.412.1913I", "2011MNRAS.414..888E", "2011MNRAS.416.1680C", "2011PASJ...63S.415T", "2012ApJ...757...15H", "2012ApJ...757...23K", "2012MNRAS.425.1320I", "2012MNRAS.425.2203T", "2012MNRAS.426..258A", "2012PASJ...64...55M", "2013ARA&A..51..207B", "2013ApJ...768...74T", "2013ApJ...772..137I", "2013ApJ...778..114T", "2013IAUS..295...74K", "2014ApJ...780...77T", "2014ApJ...785..111W", "2014MNRAS.441L...1S", "2014MNRAS.444.3357N" ]
[ "10.1088/2041-8205/788/2/L23", "10.48550/arXiv.1403.0040" ]
1403
1403.0040_arXiv.txt
\label{sec;intro} Properties of galaxies seen in the local universe strongly depend on their surrounding environments, as is well known from the morphology--density relation. Early-type galaxies are frequently observed in high-density regions such as clusters and late-type galaxies are common in low-density regions \citep{1997ApJ...490..577D}. The ATLAS$^\mathrm{3D}$ survey further demonstrates that early-type galaxies are divided into two kinds of populations based on the kinematics: fast rotators and slow rotators \citep{2011MNRAS.414..888E}. This categorization makes the trend in the morphology--density relation more prominent. Whereas fast rotators, which form the majority of early-type galaxies, appear in a wide range of environments, slow rotators reside exclusively in dense cores of mature clusters \citep{2011MNRAS.416.1680C}. Such spatial segregation and the difference in their kinematics could be related to formation processes and subsequent quenching mechanisms of star formation, which probably take place at high redshift. In theoretical models, gas-rich major mergers leading to a spin-down of remnants successfully produce simulated slow rotators and are considered to be one of the possible formation processes of slow rotators \citep{2013arXiv1311.0284N}. % Given high number densities of star-forming galaxies in proto-clusters at $z>2$ \citep[e.g.][]{2011PASJ...63S.415T}, we naturally expect a high frequency of major-merger events. What is important here is whether they are interactions between gas-dominated systems. While stellar components within galaxies are collisionless, systems comprising gas are dissipational. Therefore, gas-rich major mergers trigger an intense, dusty star formation due to shocks and an inflow of gas that has lost its angular momentum, as well as establishment of the outer profile through violent relaxation \citep{1996ApJ...464..641M}. % CO observations are critical for measuring the molecular gas mass within galaxies, $M_\mathrm{gas}$, and investigating the star-formation mode characterized by the star-formation efficiency, SFE=SFR/$M_\mathrm{gas}$, and the gas depletion timescale, $\tau_\mathrm{depl}=M_\mathrm{gas}$/SFR. CO studies at high-redshift have rapidly developed over the past years not only for very bright galaxies in the dust emission such as submillimeter galaxies (SMGs; \citealt{1998ApJ...506L...7F, 2003ApJ...597L.113N, 2005MNRAS.359.1165G, 2008ApJ...680..246T, 2010ApJ...724..233E}) but for optical/near-infrared selected galaxies such as $BzK$ galaxies \citep[e.g.][]{2010ApJ...713..686D, 2013ApJ...768...74T}. However, most of them observe high-excitation CO lines mainly using the Plateau de Bure Interferometer. This significantly affects the estimates of molecular gas mass because high-$J$ lines trace dense gas regions rather than total gas reservoirs probed by CO $J=1-0$ emission. In a cluster field at $z=0.4$, \cite{2011ApJ...730L..19G} detect the CO $J=1-0$ line from five dusty star-forming galaxies and find the environmental dependence of SFE is not seen. \cite{2012MNRAS.426..258A} also find that the SFEs of two galaxies in a proto-cluster at $z=1.5$ are comparable to that of field galaxies at similar redshift. On the other hand, \cite{2013ApJ...772..137I} discover four CO luminous galaxies across an $\sim$100 kpc region at $z=2.4$ and find that two of them have a high SFE. No conclusive result could be obtained due to a small sample size. In this paper, we report results from a deep CO $J=1-0$ observation of a proto-cluster at $z=2.5$ to search for gas-rich galaxies and see if there is any environmental effect of the star-formation mode in the formation phase of the progenitors of cluster early-type galaxies seen today. We assume the Salpeter initial mass function \citep{1955ApJ...121..161S} and cosmological parameters of H$_0$ = 70 km s$^{-1}$ Mpc$^{-1}$, $\Omega _\mathrm{M}$ = 0.3, and $\Omega _\Lambda$ = 0.7. \begin{figure} \begin{center} \includegraphics[width=0.9\linewidth]{Fig.1.eps} \caption{A 2-D map of the USS1558 proto-cluster at $z = 2.5$. Red and blue squares represent CO-detected and non-detected \ha emitters (HAEs), respectively. Coordinates are relative to the radio galaxy (a red star). The primary beam is shown by a magenta circle. Contours denote the local number densities ($\Sigma_\mathrm{5th}$) of all HAEs, in a step of 2 Mpc$^{-2}$. \label{fig;distribution}} \end{center} \end{figure}
\label{sec;discussion} The separation of $\sim$30 kpc and the velocity offset of $\sim$300 km s$^{-1}$ between ID 191 and ID 193 suggest that they are probably in the initial stage of a merger. While ID 193 is already red in the rest-frame optical and bright in the IR/radio emission, ID 191 is blue, meaning less dusty or young stellar populations. Such red-blue pairs are also observed in many SMGs \citep{2002MNRAS.337....1I}. Actually, the derived SFR of 886 $M_\odot$yr$^{-1}$ in ID193 is higher by a factor of about seven with respect to the $M_*$-SFR relation of normal HAEs ($main~sequence$; \citealt{2013ApJ...778..114T}) and is similar to that of SMGs, where extremely high star formation is thought to be driven by major mergers \citep{2008ApJ...680..246T,2010ApJ...724..233E,2012MNRAS.425.1320I}. Even the high-redshift disk galaxies fed by cold accretion of gas through cosmic filaments rarely show such extremely high SFRs \citep{2009Natur.457..451D, 2008ApJ...687...59G}. Although this classification based on colors and SFRs is not straightforward, measurements of $L_\mathrm{IR}/L'_\mathrm{CO}$ ratio are helpful for a better understanding of the star-formation mode as it reflects how molecular gas is being turned into stars. In Figure \ref{fig;COIR}, we plot the two HAEs with CO and IR/radio detections on the $L_\mathrm{IR}-L'_\mathrm{CO}$ diagram to compare them with other populations taken from the literature \citep{2012PASJ...64...55M,2007PASJ...59..117K,2001ApJ...548..681B,2012MNRAS.426..258A,1997ApJ...478..144S,2010ApJ...723.1139H,2011MNRAS.412.1913I,2013ApJ...772..137I,2010MNRAS.407.2091G}. Many previous studies at high redshift are based on the CO $J=3-2$ and $J=4-3$ line observations, leading to uncertainties in the estimates of the CO $J=1-0$ luminosities because of variations of gas excitation. Therefore, we use only objects whose CO $J=$1--0 luminosities have been directly measured. Both red and blue HAE have a higher IR luminosity compared to normal star-forming galaxies at a fixed CO luminosity and are situated close to the regime of low-redshift ULIRGs. On the other hand, some SMGs seem to follow the relation of normal star-forming galaxies. Recent analyses of galaxy morphologies in $Hubble~Space~Telescope$ images show that only 30\% of SMGs are associated with ongoing mergers \citep{2014ApJ...785..111W} while most of low-redshift ULIRGs are likely to be merger-driven starbursts \citep{2012ApJ...757...23K}. This merger fraction of SMGs is highly contentious, with several high-resolution CO studies finding it to be 100\% \citep[e.g.][]{2010ApJ...724..233E}. The starburst in ID 191 and ID193 is hard to explain by the first encounter or the final coalescence of a merger between them because the separation of 30 kpc is too large to induce violent star formation \citep{2009PASJ...61..481S}. The red HAE with the high SFR, by itself, might be in a late stage of merging where the starburst is induced by a past interaction with another galaxy. In contrast, the IR luminosity measured in the blue one is highly uncertain because of the low signal-noise ratio of 1.8$\sigma$ in the radio emission. Therefore, we do not use ID 191 to characterize the molecular gas properties. CO $J=1-0$ observations allow us to derive molecular gas masses independent of its excitation level. The difference of star-formation mode is closely related to a CO to H$_2$+He conversion factor, $\alpha_\mathrm{CO}$. Recent measurements of dust mass and dynamical mass demonstrate that $\alpha_\mathrm{CO}=3.6~M_\odot$ (K km s$^{-1}$ pc$^{-2})^{-1}$ is likely to be appropriate for normal star-forming galaxies at high-redshift \citep{2010ApJ...713..686D, 2011ApJ...740L..15M} and $\alpha_\mathrm{CO}=0.8$ is for starburst galaxies \citep{2013ARA&A..51..207B, 2013ApJ...772..137I, 1998ApJ...507..615D}. Provided that ID193 is in the starburst mode, the molecular gas mass and the gas fraction are $M_\mathrm{gas}=2.2\times10^{10}~M_\odot$ and $f_\mathrm{gas}=M_\mathrm{gas}/(M_\mathrm{gas}+M_*)\sim31\%$, respectively. Its gas depletion timescale is $\tau_\mathrm{depl}=M_\mathrm{gas}/\mathrm{SFR}\sim$ 25 Myr which is roughly consistent with that of SMGs at similar redshift \citep{2005MNRAS.359.1165G,2011MNRAS.412.1913I}. Even if using $\alpha_\mathrm{CO}=3.6$, we find that $\tau_\mathrm{depl}=110~$Myr is still shorter compared to normal star-forming galaxies (400 Myr in Salpeter IMF; \citealt{2013ApJ...768...74T}). \begin{figure}[t] \includegraphics[width=1\linewidth]{Fig.3.eps} \caption{IR vs. CO $J=1-0$ luminosities of the HAEs with CO detection along with SMGs at $z>1$ \citep[green stars:][]{2010ApJ...723.1139H,2011MNRAS.412.1913I,2013ApJ...772..137I}, ULIRGs at $z<1$ \citep[green crosses:][]{1997ApJ...478..144S}, optical/near-IR selected star-forming galaxies at $z>1$ \citep[black circles:][]{2012MNRAS.426..258A}, and local spirals \citep[black triangles:][]{2012PASJ...64...55M,2007PASJ...59..117K,2001ApJ...548..681B}. Dashed black and red lines show the best-fitting relations for normal star-forming galaxies and luminous mergers \citep{2010MNRAS.407.2091G}. \label{fig;COIR}} \end{figure} In this work, we have discovered a close pair of CO emitting HAEs, ID 191 and ID 193, in the USS1558--003 proto-cluster at $z=2.5$. Because the sum of their stellar masses becomes $M_*\sim10^{11}~M_\odot$ and subsequent star-formation activities would further increase it, such a merger system can be a good candidate of a progenitor of massive quiescent galaxies in local cluster cores. Moreover, ID 193 shows a violent star-formation activity (SFR=886 $M_\odot$yr$^{-1}$), high $L_\mathrm{IR}/L'_\mathrm{CO}$ ratio (high SFE) and red optical color. We interpret it as a star-bursting galaxy driven by a gas-rich merger with ID 191 or another galaxy. Given the identification of a rare event with the depletion timescale of $\tau_\mathrm{depl}=25$ Myr, merger events could frequently occur in the extremely dense environment as surrounding galaxies are expected to be dragged by a gravitational potential of a large halo. By dissipative processes between gas-rich galaxies, the systems that have undergone merger events could be observed as the remnants with little or no rotation (slow rotators) at $z=0$ \citep{2013arXiv1311.0284N}. This hypothesis can naturally account for the observational results that slow rotators are frequently observed in local cluster cores \citep{2011MNRAS.416.1680C}. To confirm whether gas-rich merger events are in fact common in high-density environments at $z>2$, we need a statistical sample. The combination of a wide-field \ha narrow-band imaging with a future wide-field observation of CO $J=1-0$ emission with ALMA Band-1 ($>$7 arcmin$^{2}$) will be powerful for such studies of star-formation modes in proto-clusters at high redshift. \ This paper is based on data collected at JVLA, which is operated by the National Radio Astronomy Observatory. This research has also made use of the NASA/ IPAC Infrared Science Archive, which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration. Data analysis were carried out on common use data analysis computer system at the Astronomy Data Center, ADC, of the National Astronomical Observatory of Japan. We thank the anonymous referee who gave us many useful comments, which improved the paper. K.T. and Y.K. acknowledge the support from the Japan Society for the Promotion of Science (JSPS) through JSPS research fellowships for young scientists. T.K. acknowledges the financial support in part by a Grant-in-Aid for the Scientific Research (Nos.\, 18684004, 21340045, and 24244015) by the Japanese Ministry of Education, Culture, Sports, Science and Technology. Y.T. acknowledges the support from a JSPS KAKENHI grant number 25103503.
14
3
1403.0040
Gas-rich major mergers in high-redshift proto-clusters are important events, perhaps leading to the creation of the slowly rotating remnants seen in the cores of clusters in the present day. Here, we present a deep Jansky Very Large Array observation of CO J = 1-0 emission line in a proto-cluster at z = 2.5, USS1558-003. The target field is an extremely dense region, where 20 Hα emitters (HAEs) are clustering. We have successfully detected the CO emission line from three HAEs and discovered a close pair of red and blue CO-emitting HAEs. Given their close proximity (~30 kpc), small velocity offset (~300 km s<SUP>-1</SUP>), and similar stellar masses, they could be in the early phase of a gas-rich major merger. For the red HAE, we derive a total infrared luminosity of L <SUB>IR</SUB> = 5.1 × 10<SUP>12</SUP> L <SUB>⊙</SUB> using MIPS 24 μm and radio continuum images. The L<SUB>IR</SUB>/L<SUP>\prime </SUP><SUB>CO</SUB> ratio is significantly enhanced compared to local spirals and high-redshift disks with a similar CO luminosity, which is indicative of a starburst mode. We find the gas depletion timescale is shorter than that of normal star-forming galaxies regardless of adopted CO-H<SUB>2</SUB> conversion factors. The identification of such a rare event suggests that gas-rich major mergers frequently take place in proto-clusters at z &gt; 2 and may involve the formation processes of slow rotators seen in local massive clusters.
false
[ "local massive clusters", "clusters", "CO", "local spirals", "slow rotators", "important events", "proto-clusters", "USS1558", "Gas-rich major mergers", "gas-rich major mergers", "a similar CO luminosity", "high-redshift proto-clusters", "adopted CO-H<SUB>2</SUB> conversion factors", "the CO emission line", "small velocity", "~30 kpc", "high-redshift disks", "red and blue CO-emitting HAEs", "gt", "z" ]
13.179971
7.752627
132
484283
[ "McIntosh, Scott W.", "Wang, Xin", "Leamon, Robert J.", "Davey, Alisdair R.", "Howe, Rachel", "Krista, Larisza D.", "Malanushenko, Anna V.", "Markel, Robert S.", "Cirtain, Jonathan W.", "Gurman, Joseph B.", "Pesnell, William D.", "Thompson, Michael J." ]
2014ApJ...792...12M
[ "Deciphering Solar Magnetic Activity. I. On the Relationship between the Sunspot Cycle and the Evolution of Small Magnetic Features" ]
84
[ "High Altitude Observatory, National Center for Atmospheric Research, P.O. Box 3000, Boulder, CO 80307, USA", "High Altitude Observatory, National Center for Atmospheric Research, P.O. Box 3000, Boulder, CO 80307, USA; School of Earth and Space Sciences, Peking University, Beijing 100871, China", "Department of Physics, Montana State University, Bozeman, MT 59717, USA", "Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138, USA", "School of Physics and Astronomy, University of Birmingham, Edgbaston, Birmingham, B15 2TT, UK", "Cooperative Institute for Research in Environmental Sciences, University of Colorado, Boulder, CO 80205, USA", "Department of Physics, Montana State University, Bozeman, MT 59717, USA; Lockheed-Martin Solar and Astrophysics Laboratory, 3251 Hanover Street, Org. A021S, Bldg. 252, Palo Alto, CA 94304, USA", "High Altitude Observatory, National Center for Atmospheric Research, P.O. Box 3000, Boulder, CO 80307, USA", "Marshall Space Flight Center, Code ZP13, Huntsville, AL 35812, USA", "Solar Physics Laboratory, NASA Goddard Space Flight Center, Greenbelt, MD 20771, USA", "Solar Physics Laboratory, NASA Goddard Space Flight Center, Greenbelt, MD 20771, USA", "High Altitude Observatory, National Center for Atmospheric Research, P.O. Box 3000, Boulder, CO 80307, USA" ]
[ "2014ApJ...796L..19M", "2014PhPl...21j2902L", "2015A&A...575A..63S", "2015ApJ...799...56S", "2015FrASS...2....2M", "2015NatCo...6.6491M", "2015SCPMA..58.5626L", "2015SoPh..290.3113K", "2016ApJ...830...85R", "2016ApJ...830L..27R", "2016ApJ...831..131L", "2016ApJ...833..144X", "2017A&A...601A..51I", "2017A&A...603A.134B", "2017A&A...605A..44C", "2017ApJ...845...66R", "2017ApJ...845..137G", "2017FrASS...4....4M", "2017IAUS..328...93G", "2017JASTP.162..136L", "2017NatAs...1E..86M", "2017NewA...51..135Y", "2017SoPh..292...18L", "2017SoPh..292..122K", "2017SoPh..292..181K", "2017SpWea..15.1442W", "2018A&A...609A..56C", "2018AJ....155..153K", "2018ApJS..236....5D", "2018FrASS...5...17M", "2018FrASS...5...23W", "2018FrASS...5...38S", "2018JASTP.176...15P", "2018SoPh..293....7J", "2018arXiv180307758M", "2018arXiv181202692L", "2019ApJ...873..121L", "2019ApJ...874L...4C", "2019Chaos..29f3111C", "2019LRSP...16....2M", "2019NatSR...9.2035D", "2019RvGeo..57.1129R", "2019SoPh..294...24C", "2019SoPh..294...88M", "2020AN....341..384C", "2020IJAsB..19..136A", "2020MNRAS.494.3596D", "2020RvMP...92d1001S", "2020SSRv..216..117A", "2020SoPh..295...36L", "2020SoPh..295..163M", "2020SpWea..1802109D", "2020arXiv200603111K", "2021E&SS....801223L", "2021MNRAS.502.5603R", "2021SSRv..217...15Z", "2021SoPh..296...67A", "2021SoPh..296..119E", "2021SoPh..296..174K", "2021SoPh..296..189M", "2022ApJ...930...30Z", "2022ApJS..260...38V", "2022FrASS...9.3049M", "2022FrASS...9.6670L", "2022JKPS...80..940S", "2022MNRAS.509.3717T", "2023ApJ...945...28R", "2023ApJ...945...32B", "2023ApJ...950...63M", "2023ApJ...952...42A", "2023ApJ...952..118R", "2023FrASS..1050523M", "2023FrEaS..1104191L", "2023MNRAS.523..982O", "2023arXiv230913076S", "2023hxga.book..134T", "2024A&A...683A..48H", "2024A&A...686A..90Z", "2024FrASS..1077097M", "2024MNRAS.52710440L", "2024NewA..10502103K", "2024SoPh..299...27W", "2024arXiv240212849K", "2024arXiv240517747Z" ]
[ "astronomy" ]
8
[ "Sun: activity", "Sun: atmosphere", "Sun: evolution", "Sun: general", "Sun: interior", "Sun: rotation", "sunspots", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1844AN.....21..233S", "1904MNRAS..64..747M", "1919ApJ....49..153H", "1924Natur.113..105H", "1955ApJ...121..349B", "1958Sci...127.1058B", "1959ApJ...130..364B", "1969ApJ...156....1L", "1973SoPh...32...81V", "1978ApJ...219L..55G", "1982SoPh...75..161L", "1988Natur.333..748W", "1992ASPC...27..335H", "1995SoPh..162..129S", "1995SoPh..162..291D", "1996ApJ...460.1027H", "1996Sci...272.1300T", "1999ApJ...518..508D", "1999ApJ...526..523C", "1999ApJ...527..445C", "2000Sci...287.2456H", "2003ApJ...589.1062H", "2005LRSP....2....1M", "2005LRSP....2....5S", "2005SoPh..228..285M", "2006ApJ...644L..87M", "2007ApJ...654..650M", "2007ApJ...657..486B", "2007ApJ...670.1401M", "2008ASPC..383..335A", "2009LRSP....6....1H", "2009SoPh..256...87K", "2010ApJ...725..658U", "2010LRSP....7....1H", "2010LRSP....7....3C", "2011ApJ...727....7M", "2011ApJ...731...69B", "2011ApJ...735...46R", "2011ApJ...739L..38N", "2011Natur.471...80N", "2011sswh.book...39S", "2012SoPh..275...17L", "2012SoPh..275..207S", "2013ApJ...762...73N", "2013ApJ...763...23S", "2013ApJ...765..146M", "2013ApJ...774L..29Z", "2013ApJ...778...41K", "2013SoPh..282..249T", "2013SoPh..287..239W", "2014ApJ...784L..32M", "2014SoPh..289..441N" ]
[ "10.1088/0004-637X/792/1/12", "10.48550/arXiv.1403.3071" ]
1403
1403.3071.txt
% Early investigations of the enigmatic spots on the Sun revealed that their number waxes and wanes over a period of about 11 years \citep[e.g.,][]{1844AN.....21..233S} Ð a phenomenon that became known as the sunspot (or solar) cycle \citep[e.g.,][]{1904MNRAS..64..747M}. It was subsequently found that the latitudinal distribution of sunspots, and their progression over their evolutionary cycle, followed a trail from mid solar latitudes (about $\pm$35\degr) at first appearance, through solar maximum (when their number is at its greatest) to their eventual disappearance near the equator (about $\pm$5\degr) into the relative calm of solar minimum. Following minimum, a couple of years later, the spots appear again at mid-latitudes and the progression to the equator starts afresh, defining the start of the next sunspot cycle. The pattern that sunspot locations make in this cyclic progression when latitude is plotted versus time is dubbed the ``butterfly diagram'' and has become an iconic image of the Sun's variability \citep[e.g.,][]{2010LRSP....7....1H}. Continuing this rapid pace of discovery, G.~E. Hale and colleagues subsequently identified that sunspots were locations of intense magnetic field \citep[][]{1919ApJ....49..153H} and that, in consecutive butterfly wings (sunspots in the same-hemisphere but the belonging to the next cycle), the sunspots had opposite magnetic polarities \citep[][]{1924Natur.113..105H}. Indeed, they had discovered that the sign of the prevalent magnetic field in each hemisphere of the Sun undergoes a complete period every 22 years \citep[e.g.,][]{1959ApJ...130..364B,1992ASPC...27..335H}. The radiative and particulate output of the Sun is strongly modulated by the 11-year sunspot cycle. The continued observation and cataloging of sunspots since the pioneering observations of Schwabe, Maunder, and Hale have provided us with the most striking proxy of the Sun's activity cycle which induces both violent (short-term \-- ``Space Weather'') and gradual (long-term \-- ``Space Climate'') changes in the Sun-Earth connection. The ever-increasing reliance of humanity on space-based technology has reached the point where understanding the origins and impacts of magnetic activity of our star is imperative. In the following sections we present an analysis of small ubiquitous features in the Sun's photosphere and corona that, when identified in images and tracked over a period of time, illustrate a considerably longer record of systematic magnetic evolution than sunspots. Further, the observed temporal progression appears to illustrate the latitudinal variation of oppositely signed toroidal magnetic flux systems, or ``bands'', belonging to Hale's 22-year magnetic polarity cycle. This observational finding confirms the observational work of \citet{1988Natur.333..748W} and \citet{1992ASPC...27..335H} who were first to identify the observational traces of an ``extended solar cycle'' \citep[see also][]{2013SoPh..282..249T}. Finally, we expand upon our analysis, and that of these pioneering papers, to show that the landmarks of sunspot cycle 23 can be phenomenologically described by considering the latitudinal interaction between these overlapping longer-lived bands.
% We conclude that observations covering the past seventeen years with \soho{} and \sdo{} indicate that the spatio-temporal evolution of the 11-year sunspot cycle is a result of interaction between the temporally overlapping activity bands belonging to the 22-year magnetic activity cycle. Synoptic monitoring of the solar photosphere and corona will continue in order to test this hypothesis and validate our forecast that the sunspots of solar cycle 25 will begin to appear as early as the of this decade.
14
3
1403.3071
Sunspots are a canonical marker of the Sun's internal magnetic field which flips polarity every ~22 yr. The principal variation of sunspots, an ~11 yr variation, modulates the amount of the magnetic field that pierces the solar surface and drives significant variations in our star's radiative, particulate, and eruptive output over that period. This paper presents observations from the Solar and Heliospheric Observatory and Solar Dynamics Observatory indicating that the 11 yr sunspot variation is intrinsically tied to the spatio-temporal overlap of the activity bands belonging to the 22 yr magnetic activity cycle. Using a systematic analysis of ubiquitous coronal brightpoints and the magnetic scale on which they appear to form, we show that the landmarks of sunspot cycle 23 can be explained by considering the evolution and interaction of the overlapping activity bands of the longer-scale variability.
false
[ "sunspot cycle", "significant variations", "Sunspots", "sunspots", "Solar Dynamics Observatory", "the overlapping activity bands", "Heliospheric Observatory", "interaction", "the 22 yr magnetic activity cycle", "the activity bands", "ubiquitous coronal brightpoints", "the magnetic scale", "the magnetic field", "polarity", "Solar", "the Suns internal magnetic field", "the 11 yr sunspot variation", "the longer-scale variability", "Sun", "The principal variation" ]
11.490736
15.662372
2
483201
[ "Morello, G.", "Waldmann, I. P.", "Tinetti, G.", "Peres, G.", "Micela, G.", "Howarth, I. D." ]
2014ApJ...786...22M
[ "A New Look at Spitzer Primary Transit Observations of the Exoplanet HD 189733b" ]
31
[ "Department of Physics and Astronomy, University College London, Gower Street, WC1E6BT, UK; INAF-Osservatorio Astronomico di Palermo, Piazza del Parlamento 1 I-90134, Italy;", "Department of Physics and Astronomy, University College London, Gower Street, WC1E6BT, UK", "Department of Physics and Astronomy, University College London, Gower Street, WC1E6BT, UK", "Dipartimento di Fisica, Università degli Studi di Palermo, via Archirafi I-90123, Italy.", "Dipartimento di Fisica e Chimica (previously Dipartimento di Fisica), Specola Universitaria, Università degli Studi di Palermo, Piazza del Parlamento 1 I-90123, Italy", "Department of Physics and Astronomy, University College London, Gower Street, WC1E6BT, UK" ]
[ "2014MNRAS.444.3632H", "2015ApJ...802...28C", "2015ApJ...802..107W", "2015ApJ...802..117M", "2015ApJ...808...56M", "2015ApJ...810..118K", "2015ExA....40..329T", "2016AJ....152...44I", "2016ApJ...820...86M", "2016ApJ...823..139M", "2016IJMPC..2730002S", "2017AJ....154...39D", "2017ApJS..231...13V", "2017PASP..129a4001S", "2018A&A...619A..96K", "2018AJ....156..175M", "2018ExA....46..135T", "2018MmSAI..89..332M", "2018exha.book.....P", "2020AJ....159..109M", "2020ApJ...898..119R", "2020ApJ...899...27Z", "2020MNRAS.497.5309D", "2020arXiv201015996N", "2022ApJS..258...40K", "2022ExA....53..547K", "2022MNRAS.509..289K", "2022MNRAS.512.4556S", "2023A&A...676A..54M", "2023ApJS..265....4K", "2023RASTI...2..695N" ]
[ "astronomy" ]
10
[ "methods: data analysis", "planets and satellites: atmospheres", "planets and satellites: individual: HD 189733b", "techniques: photometric", "Astrophysics - Earth and Planetary Astrophysics" ]
[ "2000ApJ...537..916S", "2000ISPL....7..197Y", "2001ApJ...553.1006B", "2002ApJ...580L.171M", "2003ApJ...585.1038S", "2007ApJ...654L..99T", "2007ApJ...668L.179E", "2007Natur.448..169T", "2008ApJ...677.1343B", "2008MNRAS.389.1383K", "2009ApJ...699..478D", "2010A&A...524A..25T", "2010ApJ...721.1861A", "2010MNRAS.403..151C", "2011A&A...526A..12D", "2011ApJ...736...12B", "2011MNRAS.418.1165H", "2012A&A...539A.140B", "2012ApJ...747...12W", "2013ApJ...766....7W" ]
[ "10.1088/0004-637X/786/1/22", "10.48550/arXiv.1403.2874" ]
1403
1403.2874_arXiv.txt
Observations of exoplanetary transits are a powerful tool to investigate the nature of planets around other stars. Transits are revealed through periodic drops in the apparent stellar brightness, due to the interposition of a planet between the star and the observer. The shape of an exoplanetary transit lightcurve depends on the geometry of the star-planet-observer system and the spatial distribution of the stellar emission at the wavelength at which observations are taken \citep{ma02}. By solving the inverse problem, it is possible to characterise fully the planet's orbit (Period, $P$; semimajor axis, $a$; inclination, $i$; eccentricity, $e$; and argument of periastron, $\omega$), and to measure its radius, $r_p$ \citep{sea03, kip08, ma02}. Knowledge of the inclination enables determination of the mass of the planet, $m_p$, if $m_p \sin{i}$ is known from radial-velocity measurements. Multiwavelength transit observations can be used to characterise the atmospheres of exoplanets, through differences in the transit depths, typically at the level of one part in $\sim 10^{4}$ in stellar flux for giant planets \citep{brown01, sea00, tin07b}. For this purpose, the diagnostic parameter is the wavelength-dependent factor $p=r_p/R_s$, i.e. the ratio between the planetary and the stellar radii (or its square, related to the transit depth). The exoplanet HD189733b is one of the most extensively studied hot Jupiters: the brightness of its star allows spectroscopic characterisation of the planet's atmosphere. The 3.6$\mu$m transit depth for the exoplanet HD189733b has been debated in the literature. Different analyses of the same dataset, including two simultaneous Spitzer/IRAC observations at 3.6$\mu$m and 5.8$\mu$m, have been used to infer the presence of water vapour in the atmosphere of HD189733b \citep{bea08, tin07}, or to reject this hypothesis \citep{des09}. Another analysis of this dataset is reported by \cite{ehr07}, but we do not comment further their results, as they were not conclusive, because of the very large error bars. \cite{des11} reported the analysis of a second Spitzer/IRAC dataset at 3.6$\mu$m using the same techniques. Their new estimates of the planet's parameters were significantly different from those reported previously by the same authors \citep{des09}; the discrepancies were attributed by the authors to variations in the star. Although stellar activity may significantly affect estimates of exoplanetary parameters from transit lightcurves \citep{bal12, ber11}, the method used to retrieve the signal of the planet also plays a critical role. The analyses mentioned above were all based on parametric corrections of the instrumental systematics, and are thus, to some degree, subjective. Recently, non-parametric methods have been proposed to decorrelate the transit signals from the astrophysical and instrumental noise, and ensure a higher degree of objectivity. \cite{wal12, wal13} suggested algorithms based on Independent Component Analysis (ICA) to extract information of an exoplanetary atmosphere from Hubble/NICMOS and Spitzer/IRS spectrophotometric datasets. In this paper we adopt a similar approach to detrend the transit signal from photometric observations by using Point Spread Functions (PSFs) covering multiple pixels on the detector. We apply this technique to re-analyse the two observations of primary transits of HD189733b recorded with Spitzer/IRAC at 3.6$\mu$m (channel 1 of IRAC) in the ``cold Spitzer'' era. We present a series of tests to assess the robustness of the method and the error bars of the parameters estimated. Critically, by comparing the results obtained for the two measurements, we discuss the level of repeatibility of transit measurements in the IR, limited by the absolute photometric accuracy of the instrument and possible stellar activity effects. We discuss the reliability of our results for orbital and stellar parameters in the light of previous multiple 8$\mu$m observations \citep{agol10}.
We have introduced a blind signal-source separation method, based on ICA, to analyse photometric data of transiting exoplanets, with a high degree of objectivity; a novel aspect is the use of pixel-lightcurves, rather than multiple observations. \\ We have applied the method to a reanalysis of two Spitzer/IRAC datasets at 3.6$\mu$m, which previous analyses found to give discrepant results, and obtained consistent transit parameters from the two observations. \\ We suggest that the large scatter of results reported in the literature arises from: \begin{itemize} \item use of arbitrary parametric methods to detrend the transit signals, neglecting the relevant uncertainties; \item correlations between parameters in the lightcurve fit. \end{itemize} We found, for observation ID 40732, values for the orbital parameters that are in excellent agreement with those found by \cite{agol10}, based on Spitzer/IRAC observations at $8 \mu m$. By applying these values to observation ID 30590, we improved the accuracy of the inferred transit depth, and strengthened the consistency between the two observations.
14
3
1403.2874
Blind source separation techniques are used to reanalyze two exoplanetary transit light curves of the exoplanet HD 189733b recorded with the IR camera IRAC on board the Spitzer Space Telescope at 3.6 μm during the "cold" era. These observations, together with observations at other IR wavelengths, are crucial to characterize the atmosphere of the planet HD 189733b. Previous analyses of the same data sets reported discrepant results, hence the necessity of the reanalyses. The method we used here is based on the Independent Component Analysis (ICA) statistical technique, which ensures a high degree of objectivity. The use of ICA to detrend single photometric observations in a self-consistent way is novel in the literature. The advantage of our reanalyses over previous work is that we do not have to make any assumptions on the structure of the unknown instrumental systematics. Such "admission of ignorance" may result in larger error bars than reported in the literature, up to a factor 1.6. This is a worthwhile tradeoff for much higher objectivity, necessary for trustworthy claims. Our main results are (1) improved and robust values of orbital and stellar parameters, (2) new measurements of the transit depths at 3.6 μm, (3) consistency between the parameters estimated from the two observations, (4) repeatability of the measurement within the photometric level of ~2 × 10<SUP>-4</SUP> in the IR, and (5) no evidence of stellar variability at the same photometric level within one year.
false
[ "other IR wavelengths", "IR", "trustworthy claims", "stellar variability", "new measurements", "single photometric observations", "observations", "the IR camera IRAC", "discrepant results", "objectivity", "larger error bars", "board", "previous work", "the same photometric level", "Blind source separation techniques", "the Spitzer Space Telescope", "~2", "orbital and stellar parameters", "the unknown instrumental systematics", "two exoplanetary transit light curves" ]
6.96768
14.466882
-1
12405223
[ "Kuijpers, Jan", "Frey, Harald U.", "Fletcher, Lyndsay" ]
2015SSRv..188....3K
[ "Electric Current Circuits in Astrophysics" ]
15
[ "Department of Astrophysics, IMAPP, Radboud University Nijmegen, Nijmegen, The Netherlands", "Space Sciences Laboratory, Berkeley, CA, USA", "SUPA School of Physics and Astronomy, University of Glasgow, Glasgow, UK" ]
[ "2015SSRv..188..211F", "2016PTEP.2016f3E01T", "2016mssf.book..231F", "2017ApJ...836..108K", "2018RvGeo..56..309G", "2019NatSR...919462G", "2019SSRv..215...51F", "2019SciA....5.8411K", "2021HPLSE...9E...2G", "2021NatCo..12.1207Z", "2021NatCo..12.2734Z", "2021NatSR..1123592G", "2021SGeo...42..757M", "2021SSRv..217...38K", "2023AcASn..64...24W" ]
[ "astronomy", "physics" ]
4
[ "Cosmic magnetism", "Electric circuits", "Radio pulsar winds", "Solar flares", "Terrestrial magnetic storms", "Astrophysics - High Energy Astrophysical Phenomena", "Astrophysics - Solar and Stellar Astrophysics", "Physics - Space Physics" ]
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[ "10.1007/s11214-014-0041-y", "10.48550/arXiv.1403.0795" ]
1403
1403.0795_arXiv.txt
\label{intro} Magnetic field structures in the cosmos occur on many scales, spanning a range of over 15 decades in spatial dimension, from extragalactic winds and jets down to the terrestrial magnetosphere. Yet in all these objects the properties of magnetic fields lead to a very similar, multi-scale, spatial and temporal structure. Magnetic fields originate in electric currents, as described by Maxwell's equations. They have energy, pressure, and tension, as quantified by their energy-momentum (stress) tensor. They exert a force on ionized matter, as expressed by the Lorentz force. Finally, their equivalent mass is small when compared to the mass-energy of ambient matter. These basic properties lead to a common appearance in a variety of objects which can be understood as follows. Since there are no magnetic charges, magnetic fields are not neutralized and they extend over large stretches of space and time. The Lorentz force which keeps ionized matter and magnetic fields together allows gravitation to anchor a magnetic field in a condensed ionized object. Since there gas and dynamic pressures dominate over the magnetic pressure, magnetic fields tend to be amplified by differential motions such as occurs in a stellar convection zone, a differentially rotating accretion disk, binary motion, or a stellar wind around a planet. Expressed in circuit language, a voltage source is set up by fluid motions which drives a current and forms a dynamo in which magnetic field is amplified at the expense of kinetic energy. Next, the small equivalent mass of the magnetic fields makes them buoyant, and as ionized matter slides easily along magnetic fields, they pop up out of the dense dynamo into their dilute environs. Whereas this central domain is dominated by fluid pressure the environs are dilute so that there the magnetic pressure dominates. The tension of the magnetic field then allows for transport of angular momentum from the central body along the field outward into a magnetized atmosphere such as a corona or magnetosphere. As more and more magnetic flux rises into the corona, and/or differential motions at the foot-points of magnetic structures continue to send electric currents and associated Poynting flux into the corona the geometries of these nearly force-free structures adapt, expand and generally lead to the appearance of thin sheets of concentrated electric currents. The same process occurs in the terrestrial magnetosphere but now the Poynting flux is going inward toward the central body. Finally, the magnetic structure in the envelope becomes unstable and the deposited energy is released in a process equivalent to an electromotor or Joule heating but now in a multitude of small `non-force-free' regions created by the currents themselves, often together in explosions, such as storms, (nano-)flares, ejected plasmoids, jets, but also in a more steady fashion, such as super-fast winds. The description of the formation and dynamics of complicated magnetic structures in terms of simplified electric current circuits in a variety of objects elucidates the fundamental physics by demanding consistency, and distinguishing cause and effect. Also, it allows for a unified answer to a number of pertinent questions: \begin{itemize} \item {\it Current Distribution}: what is the voltage source, how does the current close, which domains can be considered frozen-in, where (and when) is the effective resistivity located? \item {\it Angular Momentum Transport}: where are the balancing (decelerating and accelerating) torques? \item {\it Energy Transport}: what is the relative importance of Poynting flux versus kinetic energy flux? \item {\it Energy Conversion}: what is the nature of the effective resistivity (Lorentz force, reconnection, shocks), and what is the energy partitioning, i.e. the relative importance of gas heating, particle acceleration, bulk flow, and MHD waves? \end{itemize} For this review we have chosen to zoom in on the magnetosphere of the {\it radio pulsar}, on the {\it solar flare}, and on {\it terrestrial aurora and magnetic storms}. We will point out parallels and similarities in the dynamics of the multi-scale magnetic structures by considering the relevant electric circuits. Many of the same questions (and answers) that are addressed below are relevant to other objects as well, such as extragalactic jets, gamma ray bursts, spinning black holes, and planetary magnetospheres. Equations will be given in Gaussian units throughout to allow simple comparisons to be made while numerical values are in a variety of units reflecting their usage in the fields they come from.
We have investigated electric current systems in three very different magnetized objects in the cosmos: radio pulsar winds, the solar corona, and the terrestrial magnetosphere. Their dimensions and physical conditions differ greatly, yet their main physical effects are very similar. Essentially, this happens because, in all three cases, an electric current system is set up inside a magnetically dominated plasma by a kinematically dominated plasma, i.e. the rotating magnetic neutron star, the dynamo below the solar surface, and the solar wind. These `drivers' continue to try and increase the free energy of the electric circuit via transport of Poynting flux. However, ultimately, this results in electric current densities which in a number of small spatial domains are too large to be maintained. These domains are the locations of episodic conversion of electric current energy and the release of energized plasma in the form of ejected plasmoids, particle acceleration and heating, and magnetic turbulence. These electric circuits are largely force-free inside their magnetically dominated domains, and can be described to a first approximation within the framework of MHD. Yet we underline that it is important to consider the complete electric circuit and not just the force-free part of the magnetic field. While it is well-known that a force-free structure cannot exist on its own and has to be anchored in a non force-free bounding plasma, these boundary conditions are easy to overlook in a magnetic field picture only. When one considers an electric circuit, one automatically includes the voltage source and the resistive regions, which of course are non force-free. Although their relative volumes are minor they do determine the evolution of the system. Finally, one would like to know to what extent the end products of bulk motion, particle acceleration and heating, and the release of magnetic turbulence come about by the same physical effects in these different objects. With the proviso that we can make in situ observations inside the terrestrial magnetosphere only, while this is not possible in the corona, let alone in a pulsar wind, we conclude that possibly in all three cases one and the same process does play an important role in converting magnetic/current energy into particle acceleration. This is the occurrence of electric fields along the magnetic field whenever the local plasma conditions cannot provide for the electric current dictated by the global circuit. It is clear, however, that `common' reconnection occurs as well next to or in combination with this `Generalized Magnetic Reconnection'.
14
3
1403.0795
Cosmic magnetic structures have in common that they are anchored in a dynamo, that an external driver converts kinetic energy into internal magnetic energy, that this magnetic energy is transported as Poynting flux across the magnetically dominated structure, and that the magnetic energy is released in the form of particle acceleration, heating, bulk motion, MHD waves, and radiation. The investigation of the electric current system is particularly illuminating as to the course of events and the physics involved. We demonstrate this for the radio pulsar wind, the solar flare, and terrestrial magnetic storms.
false
[ "internal magnetic energy", "terrestrial magnetic storms", "Cosmic magnetic structures", "energy", "MHD waves", "radiation", "bulk motion", "particle acceleration", "Poynting flux", "the magnetic energy", "this magnetic energy", "MHD", "heating", "an external driver converts", "events", "the magnetically dominated structure", "the electric current system", "the radio pulsar wind", "the solar flare", "the form" ]
5.993883
3.822284
69
482758
[ "Hardegree-Ullman, E. E.", "Gudipati, M. S.", "Boogert, A. C. A.", "Lignell, H.", "Allamandola, L. J.", "Stapelfeldt, K. R.", "Werner, M." ]
2014ApJ...784..172H
[ "Laboratory Determination of the Infrared Band Strengths of Pyrene Frozen in Water Ice: Implications for the Composition of Interstellar Ices" ]
24
[ "New York Center for Astrobiology and Department of Physics, Applied Physics, and Astronomy, Rensselaer Polytechnic Institute, 110 8th Street, Troy, NY 12180, USA; Infrared Processing and Analysis Center, Mail Code 100-22, California Institute of Technology, Pasadena, CA 91125, USA;", "Jet Propulsion Laboratory, California Institute of Technology, 4800 Oak Grove Drive, Pasadena, CA 91109, USA; IPST, University of Maryland, College Park, MD 20742, USA;", "Infrared Processing and Analysis Center, Mail Code 100-22, California Institute of Technology, Pasadena, CA 91125, USA; SOFIA Science Center, USRA, NASA Ames Research Center, M.S. N232-12, Moffett Field, CA 94035, USA", "Department of Chemistry, University of California Irvine, Irvine, CA 92697-2025, USA; Division of Chemistry and Chemical Engineering, California Institute of Technology, Pasadena, CA 91125, USA", "Space Science Division, Mail Stop 245-6, NASA Ames Research Center, Moffett Field, CA 94035, USA", "NASA Goddard Space Flight Center, Exoplanets and Stellar Astrophysics Laboratory, Code 667, Greenbelt, MD 20771, USA", "Jet Propulsion Laboratory, California Institute of Technology, 4800 Oak Grove Drive, Pasadena, CA 91109, USA" ]
[ "2015ARA&A..53..541B", "2015ApJ...799...14C", "2015JChPh.143h4312C", "2015JPCA..119.2607L", "2017A&A...599A.124N", "2017ApJ...848..112D", "2017PCCP...19.8516S", "2018ApJ...864..151R", "2018PCCP...2011941M", "2018RScI...89e3111K", "2019ApJ...871..124S", "2019ApJ...875...73B", "2019FaDi..217...34B", "2019JPCA..12310694L", "2020A&A...644A..22N", "2020ApJ...902..100K", "2020ApJS..251...22M", "2021ApJ...908..239C", "2021ApJ...912...47G", "2021Carbo.177..252G", "2021PCCP...23.5817A", "2022dge..book.....W", "2023MNRAS.523.5887M", "2024JPCA..128.2772S" ]
[ "astronomy" ]
4
[ "astrochemistry", "infrared: ISM", "ISM: abundances", "methods: laboratory: molecular", "stars: protostars", "Astrophysics - Astrophysics of Galaxies" ]
[ "1971hfsa.book.....B", "1992ApJ...399..134A", "1992Natur.358...42S", "1994ApJ...433..179S", "1995ApJ...449L..69S", "1996ApJ...459..209B", "1997JPPA..104...25T", "1999ApJ...511L.115A", "1999ApJ...517..883B", "2000ApJ...544L..75B", "2001A&A...365..144D", "2001A&A...376..254K", "2002Natur.416..401B", "2002Natur.416..403M", "2003ApJ...596L.195G", "2003dge..conf.....W", "2004ApJ...607..346S", "2004ApJ...615L.177G", "2004ApJS..151...35G", "2004JPCA..108.4412G", "2005ApJS..161...53B", "2006ApJ...638..286G", "2006JCIS..294..248B", "2007ApJ...664.1264B", "2008ARA&A..46..289T", "2008ApJ...678..316B", "2008ApJ...678..985B", "2009A&A...495..837G", "2009ASPC..411..251M", "2010ApJS..189..341B", "2010MNRAS.409.1330W", "2011A&A...525A..93B", "2011EAS....46..305A", "2011IAUS..280..149P", "2012ApJ...747...13B", "2012ApJ...756L..24G", "2013ApJ...765..111K", "2013ApJ...777...73B", "2014ApJS..211....8B" ]
[ "10.1088/0004-637X/784/2/172", "10.48550/arXiv.1403.4663" ]
1403
1403.4663_arXiv.txt
Unidentified infrared bands discovered almost 40 yr ago in the interstellar medium (ISM) are now attributed to the CH and CC vibrational modes of polycyclic aromatic hydrocarbons (PAHs). They are estimated to account for 10\%$-$20\% of the total carbon reservoir in the ISM \citep{2011IAUS..280..149P}. Absorption features at 3.25 \citep{1994ApJ...433..179S,1996ApJ...459..209B,1999ApJ...517..883B}, 6.2 \citep{2001A&A...376..254K}, and 11.3 \microns\ \citep{2000ApJ...544L..75B} have been attributed to PAHs in a handful of young stellar object (YSO) spectra, while PAH emission bands are weak or absent toward embedded YSOs \citep{2009A&A...495..837G}. In molecular clouds and in the disks and envelopes of embedded YSOs, PAHs, along with other gas phase species, are expected to freeze out on the icy mantles of dust grains. In the dense ISM, ice mantles are important sites for prebiotic chemistry, where the long residence times of condensed molecules give them the best opportunity for interaction. Indeed, laboratory experiments have shown the ease of producing biomolecules via irradiation and heating of interstellar ice analogs \citep[e.g.][]{2002Natur.416..401B, 2002Natur.416..403M,2013ApJ...765..111K}. Much work has been done to understand PAH chemistry in ice matrices \citep[e.g.][]{2003ApJ...596L.195G,2004JPC...108.4412G,2004ApJ...615L.177G}. In particular, PAHs frozen in amorphous \h2o\ ice ionize quickly when irradiated by UV light. Larger PAHs remain stable up to 120~K, at which point the ice crystallizes and PAHs begin to react with the ice matrix itself \citep{2006ApJ...638..286G,2011EAS....46..305A}, while smaller PAH ions react at lower temperatures to form hydrogenated and oxygenated complex molecules \citep{2012ApJ...756L..24G}. Significant attention has also been given toward the task of quantifying PAHs in the ISM. An extensive online database at www.astrochem.org \citep{2010ApJS..189..341B,2014BoersmaArticle} contains both theoretical and experimental spectra of PAHs matrix-isolated in argon. These spectra greatly enhance our understanding of PAH signatures, but spectra of PAHs embedded in interstellar ice analogs are necessary to truly constrain PAH abundances within the dense ISM. Experiments have been carried out to determine band strengths for both neutral \citep{2005ApJS..161...53B,2011A&A...525A..93B} and ionized \citep{2007ApJ...664.1264B} pyrene frozen in water ice, but the absolute infrared (IR) band strengths reported in these cases were normalized to theoretically calculated band strengths. Sometimes, the region of interest reported was constrained between 1650 to 1000~cm$^{-1}$ (6$-$10~\microns) because the \h2o\ ice absorption there is less significant and more linear than in other regions, and none of these studies address the CH stretching mode at 3.25~\microns, which is most easily accessible with ground-based telescopes. Band strengths for naphthalene frozen in water ice have been reported as well \citep{2004ApJ...607..346S}. The aim of our work is to improve upon previous band strength estimates by recording PAH spectra in both infrared and ultraviolet-visible (UV-Vis) bands, then normalizing the results to previous, direct measurements of UV absolute band strengths \citep{Berlman1971, Dixon2005}. This paper reports our results for pyrene (C$_{16}$H$_{10}$) frozen in either \h2o\ or \d2o\ ice. Pyrene was chosen as a representative PAH molecule because it has been widely studied in the laboratory and a significant amount of laboratory data exists. The pyrene radical cation is proposed to be the carrier of some of the diffuse interstellar bands (DIBs) \citep{1992Natur.358...42S}. Additionally, pyrene is the smallest available PAH that is convenient to handle but at the same time compact with a C/H ratio of 1.9 which is closer to the larger PAH C/H ratio expected in the ISM (see Section~\ref{sec:carbon}). We included \d2o\ in our experiments because its absorption features are redshifted compared to those in \h2o, allowing easier measurement and detection of some pyrene bands which would normally be drowned out by \h2o\ features. With revised band strengths in hand, we attempted to identify the previously reported 3.25~\microns\ band \citep{1994ApJ...433..179S,1996ApJ...459..209B,1999ApJ...517..883B} in the set of YSO spectra published by \citet{2008ApJ...678..985B} and, where found, constrained PAH column densities and the contribution of PAH absorption to the 5$-$8~\microns\ absorption region. Section \ref{sec:experiment} describes the experimental setup for obtaining spectral measurements while Section \ref{sec:analysis} details our procedure for calculating absolute band strengths. Section \ref{sec:ysos} describes the astrophysical implications for quantifying PAHs in the dense ISM. Section \ref{sec:summary} summarizes our findings and suggests possible trajectories for further expansion of this work.
\label{sec:summary} PAHs account for 10\%$-$20\% of the Milky Way's carbon reservoir, yet their direct observation via absorption in the dense ISM remains difficult. Quantification of PAHs expected to be embedded in interstellar ices is hindered by the limited knowledge of the absolute band strengths, peak wavelengths, and band widths for PAHs frozen in ice matrices. Our work addresses this issue via an investigation of pyrene embedded in water ice. \begin{enumerate} \item We report the first laboratory determination of the band strength for the CH stretching mode of pyrene in water ice near 3.25~\microns. \item We report new absolute band strengths from 3$-$15~\microns\ for pyrene embedded in \h2o\ and \d2o\ ice. The band strengths reported here are roughly 50\% greater than those published by \citet{2011A&A...525A..93B}. \item We combed the data set of YSO spectra published by \citet{2008ApJ...678..985B} to look for evidence of PAH absorption at 3.25~\microns\ and used the results from our laboratory measurements to estimate PAH column densities where applicable. In our sample, neutral PAHs account for 2\%$-$9\% of the non-\h2o\ ice absorption from 5$-$8~\microns. Neutral PAHs account for 5\%$-$9\% of the carbon budget toward YSOs, depending on the particular PAH mixture present. \end{enumerate} Further laboratory work is needed to determine widths, peak positions, and band strengths for a large variety of PAH species embedded in ices. Figure~\ref{fig:lab} shows why this is critical. It is obvious from a comparison of the observed absorption of Mon R2 IRS 3 to our laboratory measurements that pyrene alone cannot explain the feature at 3052~\wavenum\ (3.25~\microns) attributed to CH stretching. Band A in our laboratory spectra is too narrow and slightly offset from the observed 3.25~\microns\ feature. Beyond 10~\microns, the CH out of plane bending features for pyrene do not align with the PAH detection at 11.3~\microns\ reported by \citet{2000ApJ...544L..75B}. In addition, the pyrene features in this region are much stronger than the reported 11.3~\microns\ feature. This apparent discrepancy supports the conclusion that PAHs around Mon R2 IRS 3 must contain a mixture of species. As illustrated in Figure~2 of \citet{1999ApJ...511L.115A}, a mixture of neutral PAHs will exhibit a strong, narrow feature at 3.25~\microns\ and much broader, weaker features beyond 10~\microns. This is because the position of the CH stretching mode is quite stable across PAH species whereas the CH out of plane bending modes, though stronger for any individual species, are much more variable in position. The result is an apparent broadening and weakening of the composite absorption feature beyond 10~\microns\ relative to the 3.25~\microns\ feature in the spectrum of a PAH mixture. Future work should also investigate the effects of other relevant ice matrices (e.g., CO$_2$) on the PAH absorption profiles. Eventually, the infrared band strengths for a variety of ionized PAHs need to be measured to compliment work with neutral species. Finally, the greatest difficulty for detecting PAH absorption features in YSO spectra is the lack of sufficient signal-to-noise due to atmospheric effects (in particular telluric CH$_4$ lines) in the ground-based data. Observational campaigns during the upcoming \textit{SOFIA} and \textit{James Webb Space Telescope} missions should effectively solve this problem.
14
3
1403.4663
Broad infrared emission features (e.g., at 3.3, 6.2, 7.7, 8.6, and 11.3 μm) from the gas phase interstellar medium have long been attributed to polycyclic aromatic hydrocarbons (PAHs). A significant portion (10%-20%) of the Milky Way's carbon reservoir is locked in PAH molecules, which makes their characterization integral to our understanding of astrochemistry. In molecular clouds and the dense envelopes and disks of young stellar objects (YSOs), PAHs are expected to be frozen in the icy mantles of dust grains where they should reveal themselves through infrared absorption. To facilitate the search for frozen interstellar PAHs, laboratory experiments were conducted to determine the positions and strengths of the bands of pyrene mixed with H<SUB>2</SUB>O and D<SUB>2</SUB>O ices. The D<SUB>2</SUB>O mixtures are used to measure pyrene bands that are masked by the strong bands of H<SUB>2</SUB>O, leading to the first laboratory determination of the band strength for the CH stretching mode of pyrene in water ice near 3.25 μm. Our infrared band strengths were normalized to experimentally determined ultraviolet band strengths, and we find that they are generally ~50% larger than those reported by Bouwman et al. based on theoretical strengths. These improved band strengths were used to reexamine YSO spectra published by Boogert et al. to estimate the contribution of frozen PAHs to absorption in the 5-8 μm spectral region, taking into account the strength of the 3.25 μm CH stretching mode. It is found that frozen neutral PAHs contain 5%-9% of the cosmic carbon budget and account for 2%-9% of the unidentified absorption in the 5-8 μm region.
false
[ "ultraviolet band strengths", "theoretical strengths", "pyrene bands", "strengths", "water ice", "frozen interstellar PAHs", "infrared absorption", "Boogert et al", "frozen PAHs", "Bouwman et al", "frozen neutral PAHs", "pyrene", "Our infrared band strengths", "polycyclic aromatic hydrocarbons", "These improved band strengths", "account", "interstellar medium", "the band strength", "absorption", "PAHs" ]
12.11957
11.102954
-1
483726
[ "van Eysden, C. A." ]
2014ApJ...789..142V
[ "Short-period Pulsar Oscillations Following a Glitch" ]
9
[ "NORDITA, Roslagstullsbacken 23, SE-10691 Stockholm, Sweden" ]
[ "2015IJMPD..2430008H", "2015JFM...783..251V", "2015MNRAS.454.4400N", "2016MNRAS.461L..77H", "2017PhRvD..95b4022S", "2018ASSL..457..401H", "2018ApJ...865...60V", "2018PASA...35...20K", "2022NewA...9001655P" ]
[ "astronomy" ]
4
[ "magnetohydrodynamics: MHD", "stars: neutron", "stars: rotation", "Astrophysics - Solar and Stellar Astrophysics", "Astrophysics - High Energy Astrophysical Phenomena" ]
[ "1963JFM....17..385G", "1967JFM....28..463P", "1969JFM....36..289W", "1969Natur.224..673B", "1971JFM....50..609L", "1971NuPhA.175..225B", "1975Natur.256...25A", "1976ApJ...206..218F", "1977PhRvD..16..275E", "1979ApJ...228..257E", "1979ApJ...233..711E", "1983bhwd.book.....S", "1984ApJ...282..533A", "1984ApL....24...91L", "1987ApJ...314..234C", "1989PThPh..82..945T", "1990Natur.346..822M", "1991ApJ...380..530M", "1993JLTP...92...77R", "1995ApJ...447..324S", "1996JFM...312..327A", "1998AIPC..443....3J", "1998ApJ...492..267R", "1998MNRAS.296..903M", "2001ApJ...548..447W", "2004ApJ...613.1157L", "2005NuPhA.763..212A", "2006MNRAS.368L..35L", "2006MNRAS.371L..74G", "2007MNRAS.377..159L", "2007PhR...442..109L", "2008CQGra..25v5020V", "2008MNRAS.385L...5S", "2008MNRAS.390..175W", "2008PhRvD..78f3006S", "2009ApJ...700.1524M", "2009MNRAS.396..894A", "2009MNRAS.396.1441C", "2009MNRAS.397.1607C", "2010MNRAS.402.1027H", "2010MNRAS.405..318L", "2010MNRAS.409.1253V", "2010Sci...329..408L", "2011MNRAS.410..805G", "2011MNRAS.410.1036V", "2011MNRAS.410L..37G", "2011MNRAS.412.1730L", "2011MNRAS.413...47P", "2011MNRAS.414.1679E", "2011MNRAS.414.3014C", "2011MNRAS.415.1611W", "2012ApJ...761...32M", "2012MNRAS.420..658H", "2012MNRAS.420.3035V", "2012PhRvB..85j4503W", "2013JFM...729..180V", "2013MNRAS.428.1911W", "2013PhRvL.111u1102G", "2014JFM...744...89V", "2014MNRAS.438..156P" ]
[ "10.1088/0004-637X/789/2/142", "10.48550/arXiv.1403.1046" ]
1403
1403.1046_arXiv.txt
In dynamic models of pulsars, it is almost ubiquitously assumed that the strong magnetic field couples the proton-electron plasma in the outer core to the crust on short time scales, so that the two components are rigidly locked together \citep{lod84,tak89,alp84,sed95b,pas11}. Originally proposed by \citet{bay69a}, this assumption has been validated for the slow spin-down of the pulsar frequency produced by magnetic dipole braking \citep{eas79c}, as well as for an impulsive change in the pulsar frequency induced by a glitch \citep{eas79a}. In the latter case, co-rotation is established by an Ekman-like process, driven by the Coriolis force in the magneto-inertial boundary layers that form immediately after the glitch. The spin-up time is approximately two seconds for typical pulsar parameters in an ideal magnetohydrodynamic plasma, and estimated to be a factor of thirty shorter if the protons are a type II superconductor. Questions arise, however, when one considers the differences between the classical spin up of a viscous fluid \citep{gre63} and that in an ideal magnetized plasma. In particular, co-rotation in classical spin up is achieved via a dissipative process (viscosity), whereas ideal magnetohydrodynamics is dissipation-less. In both cases, a shear gradient is established immediately after the impulsive spin up of the container in the plasma abutting the boundary. In ideal magnetohydrodynamics, however, the corresponding shear gradient in the magnetic field stores potential energy in the form of magnetic field line tension, which provides the restoring force for Alfv\'en waves that subsequently propagate through the plasma. Because the plasma is dissipation-less, these waves are continually reflected internally by the container walls, generating torsional oscillations between the container and plasma. Oscillations of this nature have received significant attention recently in attempts to explain the quasi-periodic oscillation modes present in magnetar giant flares \citep{gla06b,lev07,sot08,col09,cer09,col11,gab11,van11c}, and earlier in the context of neutron star precession \citep{lev04}. An argument for oscillations can be made purely on the grounds of conservation of angular momentum and energy; a final state of uniform co-rotation of the fluid and its container is not consistent with energy conservation and can {\it only} be achieved if there is dissipation. The question arises then, why does the result of \citet{eas79a} suggest a final state of co-rotation, and what happened to the torsional oscillations of the fluid and container? In this paper, we re-visit the problem of the spin-up of a rapidly rotating magnetized plasma after the impulsive acceleration of its container. We consider the same cylindrical geometry studied by \citet{eas79a}, but solve self-consistently for the motion of the plasma {\it and} its container. We derive a general solution and compare our results with those of \citet{eas79a}, who considered only the limit in which the ratio of the rotational period to the Alfv\'en crossing time approaches zero. We find that torsional oscillations commence after an Alfv\'en crossing time, $\sim 30\,{\rm sec}$ in a typical neutron star, and are not seen in \citet{eas79a} because the Alfv\'en crossing time is assumed to be infinitely long. The crust oscillation amplitude also diminishes with increasing Alfv\'en crossing time because the spectrum of magneto-inertial oscillations in the core approaches a continuum, where the crust motion is damped by resonant absorption in the core analogous to magnetar oscillations \citep{lev06,lev07}. There are two primary motivations for this study. First, to identify and characterize the oscillatory behavior predicted above. In the case of magnetars, extensive effort has been devoted into determining the coupled crust-core oscillation spectrum for a spherical star with realistic magnetic field configurations using numerical codes \citep{col11,gab11,van12} and analytic approaches \citep{lev07,van11c}. The effects of a superfluid core have also been investigated \citep{and09,gab13}. However, in the present paper we are concerned with rotation-powered pulsars, where, in contrast with magnetars, the rotational inertia dominates the magnetic forces. Our second motivation is to determine the motion of the crust in a self-consistent manner, so that a direct comparison with radio timing data can be made, as in \citet{van10}. The identification and characterization of oscillations in glitch recovery could provide another dimension for using glitches to constrain properties of the pulsar interior. The paper is structured as follows: The assumptions of the model and the governing equations are presented in \S\ref{sec2a}, and the solution to the system is presented in \S\ref{sec2b}. In \S\ref{sec3a} the generic properties of the solution are explored, while in \S\ref{sec3b} the regime relevant to neutron stars is presented and the observational consequences discussed. The conclusions are presented in \S\ref{sec4}.
\label{sec4} We have shown that co-rotation of the crust and plasma in the outer core of the pulsar following a glitch is not brought about by the magnetic field. Rather, the glitch excites magneto-inertial waves that propagate through the plasma, generating torsional oscillation modes of the crust as they reflect internally. The oscillations have a period of $1-20\,{\rm s}$ and decay over a timescale of $15-30\, {\rm mins}$ due to the viscosity of the electrons in the core. The toy model presented here is unlikely to represent a realistic oscillation spectrum following a glitch, however, it sheds light on some qualitative features of the problem. First, co-rotation between the plasma and crust of a neutron star cannot be achieved by the magnetic field following a glitch. This is a violation of conservation of energy; the system oscillates persistently until it is damped by electron viscosity. Second, an Ekman pumping mechanism, first identified by \citet{eas79a} is present at short times, which spins up the plasma in the interior. However, following the spin-up, Alf\'en waves excited by the glitch continue to propagate through the plasma, exciting oscillations of the crust after an Alfv\'en crossing time. Third, for rapidly rotating stars (where rotation energy dominates the magnetic energy, i.e., $\xi \rightarrow 0$), the spectrum of magneto-inertial oscillations in the core becomes a continuum and the oscillations of the container undergo Landau damping. This in analogous to the resonant absorption mechanism identified in magnetars, where the oscillation modes in the core conspire so that the net back-reaction on the crust vanishes. To determine a realistic spectrum for a neutron star following a glitch the use of a numerical code such as those used to model magnetar quasi-period oscillations is most likely required. These codes have achieved success in reproducing the observed QPO frequencies for non-rotating stars using relativistic ideal MHD, and have recently incorporated the essential physics of superfluidity, superconductivity and stratification, although these studies are still in their infancy. It would be interesting to incorporate magnetic fields into models for rotating stars and glitch recovery, as it is possible that there is physics occurring that has not been resolved by existing radio telescopes, but which may be observable by future radio telescope arrays such as LOFAR and the SKA, gravitational wave detectors and x-ray observatories. The discovery and identification of such a spectrum following a glitch would be extremely useful for constraining the physics of neutron star interiors.
14
3
1403.1046
Following a glitch, the crust and magnetized plasma in the outer core of a neutron star are believed to rapidly establish a state of co-rotation within a few seconds by process analogous to classical Ekman pumping. However, in ideal magnetohydrodynamics, a final state of co-rotation is inconsistent with conservation of energy of the system. We demonstrate that, after the Ekman-like spin up is completed, magneto-inertial waves continue to propagate throughout the star, exciting torsional oscillations in the crust and plasma. The crust oscillation is irregular and quasi-periodic, with a dominant frequency of the order of seconds. Crust oscillations commence after an Alfvén crossing time, approximately half a minute at the magnetic pole, and are subsequently damped by the electron viscosity over approximately an hour. In rapidly rotating stars, the magneto-inertial spectrum in the core approaches a continuum, and crust oscillations are damped by resonant absorption analogous to quasi-periodic oscillations in magnetars. The oscillations predicted are unlikely to be observed in timing data from existing radio telescopes, but may be visible to next generation telescope arrays.
false
[ "quasi-periodic oscillations", "crust oscillations", "next generation telescope", "classical Ekman pumping", "Crust oscillations commence", "exciting torsional oscillations", "existing radio telescopes", "seconds", "Ekman", "magnetized plasma", "process", "resonant absorption", "-", "rotation", "co", "plasma", "The crust oscillation", "energy", "magneto-inertial waves", "conservation" ]
5.196247
3.080106
69
471971
[ "Fryer, Chris L.", "Even, Wesley", "Grefenstette, Brian W.", "Wong, Tsing-Wai" ]
2014AIPA....4d1014F
[ "Observational constraints of stellar collapse: Diagnostic probes of nature's extreme matter experiment" ]
5
[ "CCS Division, Los Alamos National Laboratory, Los Alamos, NM 87545, USA", "CCS Division, Los Alamos National Laboratory, Los Alamos, NM 87545, USA", "Space Radiation Lab, California Institute of Technology, Pasadena, CA 91125, USA", "Center for Interdisciplinary Exploration and Research in Astrophysics (CIERA) &amp; Department of Physics and Astronomy, Northwestern University, 2145 Sheridan Road, Evanston, IL 60208, USA; Harvard-Smithsonian Center for Astrophysics, 60 Garden St., Cambridge, MA 02138, USA" ]
[ "2014AIPA....4d0901B", "2016ApJ...819..108B", "2016IAUFM..29B.206F", "2017ApJ...834...19G", "2018JCoPh.359...20C" ]
[ "astronomy", "physics" ]
5
[ "Astrophysics - High Energy Astrophysical Phenomena" ]
[ "1978ApJ...221L..29B", "1978MNRAS.182..147B", "1978MNRAS.182..443B", "1982ApJ...258..790C", "1982ApJ...259..302C", "1985AJ.....90.2089B", "1987PhRvL..59..736B", "1989GeCoA..53..197A", "1990ApJ...356..272W", "1991AJ....102..303R", "1991ApJ...371..621F", "1991SSRv...58..259J", "1992AJ....103.1987B", "1992ApJ...391..246P", "1992ApJ...395..202W", "1992ApJ...399..656M", "1993A&A...273..106B", "1993A&A...279..447M", "1993ApJ...417..713H", "1993ApJ...418..386L", "1993SSRv...66..439P", "1994A&A...286..857T", "1994ApJ...433..229W", "1994ApJ...435..339H", "1994PhRvL..73.1878F", "1995ApJ...444..312C", "1995ApJ...450..830B", "1995ApJ...453..792F", "1996ApJ...471..254R", "1996ApJ...471..331Q", "1996ApJ...472..257B", "1996ApJ...472..440M", "1997AIPC..412..160F", "1998ApJ...499..367B", "1998PhRvC..58.3696M", "1998PhRvD..58a3009R", "1999A&A...351..459C", "1999ApJ...512..288T", "1999ApJ...517..193N", "1999ApJ...522..413F", "2000ApJ...540..292E", "2000ApJ...541.1033F", "2001ApJ...562..887T", "2002ApJ...565..430F", "2002ApJ...567..515Q", "2002ApJ...567..532H", "2002ApJ...574L..65F", "2003APh....18..433F", "2003ApJ...584..971B", "2003ApJ...591..288H", "2003ApJ...593L..23C", "2004A&A...416..997A", "2004ApJ...601..391F", "2004ApJ...606..116F", "2004ApJ...609..288F", "2004ApJ...611.1005G", "2005A&A...437L...7A", "2005ARA&A..43..435H", "2005ApJ...625..324W", "2005ApJ...628..914S", "2005PASP..117..773V", "2005Sci...308.1604K", "2006ApJ...640..891Y", "2006ApJ...642..401B", "2006ApJ...645.1331M", "2006ApJ...646L.131F", "2006ApJ...650.1028F", "2006NewAR..50..492F", "2006PhRvD..74l3004D", "2006RvMP...78..755R", "2007ApJ...654.1006F", "2007ApJ...659..561M", "2007ApJ...659.1438F", "2007ApJ...661.1013L", "2007ApJ...662L..55F", "2007ApJ...662L.107V", "2007ApJ...664.1102K", "2007ApJ...670..747K", "2007JCAP...12..010F", "2007PASP..119.1211F", "2007PhRvD..76h1301R", "2007PhRvD..76l5008R", "2008ARA&A..46...89R", "2008JPhCS.125a2007K", "2008MNRAS.391L...5C", "2008PhRvD..78h5012E", "2009ARA&A..47...63S", "2009ApJ...697.1057F", "2009ApJ...699..409F", "2009ApJ...706.1174E", "2009ApJ...707.1173M", "2009JPhG...36k3201D", "2009MNRAS.395.1409S", "2009NIMPA.600..568A", "2009PhRvL.103e1105D", "2010ARNPS..60..569D", "2010ApJ...725.1918O", "2010ApJS..191...66M", "2010CQGra..27q3001A", "2010Natur.467.1081D", "2010PrPNP..64..273K", "2011ApJ...735..120Y", "2011ApJ...741..103F", "2011ApJ...741L..28V", "2011JPhG...38c5201D", "2011LRR....14....1F", "2011MNRAS.417.1417F", "2011NuPhS.221..173M", "2011NuPhS.221..193M", "2011PhRvD..83e2010A", "2011PhRvL.106i1101D", "2012APh....35..685W", "2012ApJ...746..130H", "2012ApJ...747...23S", "2012ApJ...749...91F", "2012ApJ...751....4K", "2012ApJ...756..131V", "2012ApJ...757...36K", "2012ApJ...757...55O", "2012ApJ...757...91B", "2012ApJ...759...52D", "2012ApJ...759..107K", "2012IAUS..279..183A", "2012PhRvC..86f5803R", "2012PhRvL.108f1103R", "2012PhRvL.108k2502A", "2012PhRvL.109y1104M", "2013A&A...552A.102T", "2013A&A...552A.105V", "2013A&A...558L...1G", "2013AJ....145...13A", "2013AJ....146..132L", "2013APh....43...71A", "2013ApJ...766L..14H", "2013ApJ...767L...9G", "2013ApJ...768..148M", "2013ApJ...770...66H", "2013ApJ...770...93A", "2013ApJ...770..103H", "2013ApJ...772...30D", "2013ApJ...772L..32V", "2013ApJ...773...11H", "2013ApJ...775L...7C", "2013ApJ...779L...9H", "2013ApJS..204...16F", "2013PhRvC..88b5501A", "2013PhRvD..88d3007A", "2013SSRv..178..599B", "2013arXiv1305.4137E", "2013arXiv1307.7419B", "2013arXiv1310.4340D", "2014ApJ...781...40P", "2014ApJ...781...42O", "2014JCAP...11..030B", "2014MNRAS.438.1577M", "2014Natur.506..339G", "2015ApJ...805...98B" ]
[ "10.1063/1.4870404", "10.48550/arXiv.1403.3619" ]
1403
1403.3619_arXiv.txt
Introduction: Supernovae as Extreme-Matter Experiments} Core-Collapse Supernovae are among the most powerful explosions in the universe, produced when the core of a massive star is no longer able to support itself and collapses. The collapse continues until the core reaches nuclear densities at which point nuclear forces and neutron degeneracy pressure halt the collapse. The compression produces densities in excess of nuclear densities and temperatures above 10\,MeV. Although laboratory experiments such as PREx \cite{abrahamyan12} might allow scientists to probe properties of the neutron, core-collapse supernova have the potential to probe the broader properties of nuclear matter at densities above nuclear. Nuclear physics is at the heart of the core-collapse supernova problem. The behavior of matter at nuclear densities determines the depth of the bounce which, in turn, can alter the strength of the bounce shock causing a prompt explosion\cite{baron87}. Although this prompt explosion mechanism is no longer favored, the behavior of matter at nuclear densities still plays a role in the explosion, defining the neutrino luminosity and spectra, determining the maximum compression of the core and altering the extent of the bounce shock. Similarly, uncertainties in the nuclear reaction rates alter both the explosive and r-process yields in supernovae. In this paper, we review core-collapse supernovae from the point-of-view of a nuclear physics experiment. As with many laboratory experiments, nuclear physics is only one piece of the core-collapse supernova problem (see review by Hix et al. in this issue), and all this physics must be understood and/or its uncertainties limited if we are to study any piece of the physics. Fortunately, there are a broad set of diagnostics probing core-collapse supernovae ranging from direct probes of the core limited to only the most nearby supernovae (neutrinos and gravitational waves) to integrated probes that include many objects (nucleosynthetic yields or the mass distribution of compact remnants). But to truly take advantage of this data, theoretical models are critical. New theory models of the progenitors, the supernova engine and its subsequent explosions have been advancing the physics and strengthening our ability to use these diagnostics to probe the physics behind core-collapse supernovae. This approach is not so different than many laboratory experiments. Although experimentalists strive to design an experiment to directly observe the intended physics, most experiments are driven by a wide range of physics and theoretical models are often needed to reduce the uncertainties and interpret the results. Before we discuss the diagnostics in detail, let's review the basic supernova engine mechanism and the stages at which the diagnostics help us pin down this engine. We will revisit aspects of this explosion mechanism throughout the paper as they pertain to specific diagnostic probes. The supernova progenitor, a star more massive than $\sim 8M_\odot$ progresses through a series of burning phases, building up a dense core. For most pre-supernova progenitors, an iron core is produced that continues to grow until it can no longer support itself, collapsing to form a neutron star. Some super-asymptotic giant branch stars can collapse without forming an iron core, the so-called electron capture supernovae\cite{herwig05}. Supernova progenitors can be observed by searching for the absence of progenitors in survey data after the supernova fades and the list of these observed stars has increased steadily over the past decade (see Sec.~\ref{sec:progenitor}). The iron core is supported by thermal and electron degeneracy pressure. As the mass increases, the core contracts until electron capture and the dissociation of iron atoms remove both these pressure sources, causing the core to collapse. The compression accelerates the capture and dissociation, leading quickly to a runaway collapse. The collapse continues until the core reaches nuclear densities where neutron degeneracy pressure and nuclear forces halt the collapse, causing the imploding star to bounce. The bounce shock stalls and the current "standard" picture of supernovae argues that convective instabilities (e.g. Rayleigh-Taylor\cite{herant94}, advective-vortical instabilities\cite{blondin03}, or advective-acoustic\cite{foglizzo07}) revive the shock, driving an explosion. Neutrinos emitted in the core increase during the collapse but, in the extreme densities at bounce can be trapped in the core. As the bounce shock expands, the neutrinos leak out, producing the strong burst of neutrinos and a near-direct probe of the behavior of matter at nuclear densities. The neutrino signal persists as the neutron star cools and can be further enhanced by late-time accretion (caused as material initially swept up in the shock is unable to escape the gravitational pull of the neutron star) onto the supernova. Gravitational waves also probe this core region during the bounce and convective engine time frame, but are more sensitive mass motions caused by rotation and convection. Elements are synthesized in a variety of processes: explosive nucleosynthesis as the shock moves through the silicon layer in the star causing fusion up through iron peak elements, r-process elements produced in winds from the cooling proto-neutron star or from secondary ejecta caused by accretion onto the newly formed neutron star. Measuring these nucleosynthetic yields in stars or asteroids provides another clue into the supernova engine (albeit in an integrated sense). As the supernova shock breaks out of the star, photons trapped in the shock are also able to leak out. This shock breakout marks the first optical emission from the supernova explosion. The shock continues to expand, producing the bright emission whose details characterize the broad range of Ib,Ic, IIp, IIn, IIl, IIb observed supernova classes. Observations of the breakout probes the photosphere of the star, whereas the peak and late-time light curve probe both characteristics of the star (photosphere, circumstellar medium, sometimes mass) as well as the yield of $^{56}$Ni. Supernova remnants, both the compact collapsed cores (neutron stars and black holes) and the ejecta remnants provide further diagnostics of the supernova explosion. Hundreds of years after the explosion, the supernova ejecta is still visible as a remnant. The supernova ejecta-remnants provide yet further clues to the explosion asymmetry and nucleosynthetic yields. The mass distribution and spin of compact remnants (the black hole or neutron star formed during collapse) can also provide clues in understanding the supernova engine. In this review, we study all of these diagnostics. First, we focus on the most direct diagnostics of the nuclear physics: neutrinos (Sec.\ref{sec:neutrinos}) that probe the collapsed core with the neutrinosphere touching the surface of the newly-formed neutron star and gravitational waves (Sec.\ref{sec:GW}) that probe the matter motion near the proto-neutron star. We then turn to the more integrated diagnostics. These diagnostics are much more common, but analyzing the data requires much more complex models with integrated physics. To constrain these diagnostics, we must understand the progenitors (Sec.~\ref{sec:progenitor}) and obtain as much data as possible on the supernova outburst itself including light curves and spectra (Sec.~\ref{sec:SNLC}), the supernova ejecta (Sec.~\ref{sec:ejectaremnant}), mass distributions of the compact remnants (Sec.~\ref{sec:compactremnant}), and the integrated nucleosynthetic yield estimates (Sec.~\ref{sec:nucobs}). As an example of the strength of these combined diagnostics, we review the study of r-process elements from core-collapse supernovae (Sec.~\ref{sec:postexp}).
Supernovae are indeed Nature's laboratories for matter in extreme conditions. Traditional studies of this extreme physics have focused on the neutrino and gravitational wave diagnostics of supernovae. But nearly all the data we have on supernovae used different diagnostics. In this review, we have discussed the range of diagnostics: neutrinos, gravitational waves, progenitor studies, the supernova outbursts (a wide range of photon energies), compact remnant mass distributions, ejecta remnant observations, and integrated nucleosynthetic yield constraints from low-metallicity stars and meteoritics. All of these diagnostics have strengths and weaknesses, but taken together, we are at the stage of truly constraining nuclear and particle physics with our supernova ``experiment''. Some of these diagnostics require considerable theory work to interpret and any study of supernovae as physics laboratories will require a strong theoretical understanding. The wealth of data will be able to both help us understand the uncertainties in this theoretical understanding as well as provide constraints on our errors in studying nuclear physics. Supernova studies are at a critical point where new surveys and new observational experiments will soon drastically increase the amount of data. Along with advances in our theoretical understanding, we are on the cusp of a revolution in this science and it is likely that the next decade will see large breakthroughs in this science.
14
3
1403.3619
Supernovae are Nature's high-energy, high density laboratory experiments, reaching densities in excess of nuclear densities and temperatures above 10 MeV. Astronomers have built up a suite of diagnostics to study these supernovae. If we can utilize these diagnostics, and tie them together with a theoretical understanding of supernova physics, we can use these cosmic explosions to study the nature of matter at these extreme densities and temperatures. Capitalizing on these diagnostics will require understanding a wide range of additional physics. Here we review the diagnostics and the physics neeeded to use them to learn about the supernova engine, and ultimate nuclear physics.
false
[ "nuclear densities", "ultimate nuclear physics", "densities", "additional physics", "temperatures", "supernova physics", "excess", "these extreme densities", "matter", "Natures high-energy, high density laboratory experiments", "diagnostics", "the supernova engine", "these cosmic explosions", "Nature", "a wide range", "a theoretical understanding", "Astronomers", "Supernovae", "the nature", "the physics" ]
3.987303
3.228008
55
527118
[ "Rybka, Gray" ]
2014arXiv1403.6720R
[ "Improving Dark Matter Axion Searches with Active Resonators" ]
2
[ "-" ]
[ "2018arXiv180301627C", "2021JCAP...12..033C" ]
[ "astronomy", "physics" ]
4
[ "Physics - Instrumentation and Detectors", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
[ "1977PhRvD..16.1791P", "1977PhRvL..38.1440P", "1978PhRvL..40..223W", "1978PhRvL..40..279W", "1983PhLB..120..127P", "1983PhLB..120..133A", "1983PhRvL..50..925I", "1983PhRvL..51.1415S", "1985PhRvD..32.2988S", "1990PhR...197...67T", "1998PhRvL..80.2043H", "2000NIMPA.444..569P", "2010PhRvL.104d1301A", "2013RScI...84h4703J" ]
[ "10.48550/arXiv.1403.6720" ]
1403
1403.6720_arXiv.txt
14
3
1403.6720
Axions are a well motivated candidate for dark matter. The most sensitive experiments searching for dark matter axions rely on the coupling of axions to the electromagnetic resonances of a microwave cavity immersed in a strong magnetic field. The sensitivity of the experiment is proportional to the Q of the resonance that is coupled to axions. To date, the resonators used in axion searches have all been passive, with Qs limited by power loss in the cavity walls. I propose the use of active feedback resonators to increase the Q of microwave cavity axion dark matter experiments by several orders of magnitude. This should allow experiments to significantly increase the rate at which they can test potential axion masses and couplings.
false
[ "microwave cavity axion", "dark matter axions", "axion searches", "Axions", "axions", "potential axion masses", "dark matter", "couplings", "power loss", "several orders", "magnitude", "experiments", "active feedback resonators", "a strong magnetic field", "a microwave cavity", "the cavity walls", "Qs", "the electromagnetic resonances", "the coupling", "the resonance" ]
9.168475
-2.323721
65
12359329
[ "Morel, Thierry" ]
2015ASSP...39...73M
[ "Spectroscopic Constraints for Low-Mass Asteroseismic Targets" ]
2
[ "Institut d'Astrophysique et de Géophysique, Université de Liège" ]
[ "2014A&A...564A.119M", "2017AN....338..644M" ]
[ "astronomy" ]
4
[ "Physics", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "2000A&A...359..743A", "2002A&A...389..345B", "2006A&A...448..341G", "2007A&A...469..687C", "2007AJ....133.2464L", "2010A&A...522A..10C", "2010A&A...523A..91G", "2010MNRAS.405.1907B", "2011A&A...528A.121B", "2011JPhCS.328a2010M", "2012A&A...537A.111C", "2012A&A...543A.160T", "2012ASPC..458..393F", "2012ApJ...757..161T", "2012ApJ...760...32H", "2012MNRAS.423..122B", "2012MNRAS.427...50L", "2012Msngr.147...25G", "2013A&A...549A..12M", "2013A&A...549A..14K", "2013A&A...552A..42M", "2013A&A...559A.102D", "2013AJ....146..133M", "2013ARA&A..51..353C", "2013ApJ...765L..41S", "2013ApJ...766..118M", "2013ApJ...767..127H", "2013MNRAS.429..423M", "2013MNRAS.434.1422M", "2014A&A...562A..71B", "2014A&A...564A.119M" ]
[ "10.1007/978-3-319-10993-0_8", "10.48550/arXiv.1403.5103" ]
1403
1403.5103_arXiv.txt
\label{sect_introduction} The great potential of asteroseismology to address some unresolved issues in stellar physics and even, as was discussed during this meeting, to study the stellar populations making up our Galaxy cannot be overstated. Yet these expectations cannot be completely met if some fundamental quantities that are not encoded in seismic data are not accurately known \citep[e.g.,][]{creevey12}. For this reason, by providing the effective temperature and chemical composition (but also other important information such as the $v\sin i$ or the binary status), a traditional field such as stellar spectroscopy will still play an important role in the future for the study of seismic targets. Conversely, asteroseismology can provide the fundamental quantities (e.g., mass, age, evolutionary status in the case of red giants) that are needed to best interpret the abundance data. These two fields are therefore closely connected and can greatly benefit from each other. The large discrepancies between the $\log g$ and [Fe/H] values derived from spectroscopy and those in the {\it Kepler} Input Catalog \citep{bruntt12,thygesen12} illustrate the clear superiority of spectroscopic techniques over photometric ones for the estimation of these two parameters. Determining accurate temperatures from photometric indices is also challenging in the presence of a significant (and patchy) reddening (e.g., for some CoRoT fields that lie close to the Galactic plane).
\label{sect_samples} Numerous spectroscopic analyses of individual seismic targets have been conducted during the last few years \citep[e.g.,][]{mathur13,morel13}. However, we will restrict ourselves here to discussing the results of studies dealing with a sizeable number of stars observed by either the CoRoT or the {\it Kepler} space missions. The CoRoT satellite operated either through the seismology (observations of a limited number of bright stars in the context of seismic studies) or the exoplanet (observations of numerous faint stars to detect planetary transits) channel. The parameters of a large number of stars in various evolutionary stages in the CoRoT exofields have been determined using an automated pipeline by \citet{gazzano10}, while a more comprehensive analysis of 19 red giants in the seismology fields has been presented by \citet{morel14}.\footnote{Note that the sample of \citet{morel14} discussed in the following contains a few stars which were eventually not observed by the satellite, as well as a number of benchmark stars used for validation purposes.} In the latter case, a standard analysis is employed that imposes excitation and ionisation equilibrium of iron based on the equivalent widths of a set of Fe I and Fe II lines. On the other hand, a study of dwarfs and giants in the {\it Kepler} field has been performed by \citet{bruntt12} and \citet{thygesen12}, respectively (the latter study superseding that of \citealt{bruntt11}). In both cases, the analysis relied on the spectral-synthesis software package {\tt VWA} \citep[see, e.g.,][]{bruntt02}. Table~\ref{tab_uncertainties} gives for all the studies mentioned above the uncertainties associated to the determination of the parameters. Based on the (sometimes rather scanty) information provided in these papers, it may be concluded that these figures are claimed to be representative of the {\it accuracy} of the results. Although these measurements also suffer from limitations (e.g., calibration issues, angular diameter corrections, reddening), the satisfactory agreement with the less model-dependent estimates provided by interferometry for stars at near-solar metallicities \citep[e.g.,][]{bruntt10,huber12,morel14} suggests that the values quoted in Table~\ref{tab_uncertainties} for $T_{\rm eff}$ are reasonable in this metallicity regime (however, this may not be true for metal-poor stars where non-LTE and 3D effects become important; \citealt{lind12,dobrovolskas13}). Much more extensive and stringent tests can be expected in the future thanks to the advent of new long-baseline interferometric facilities. A comparison for a subset of {\it Kepler} targets between the parameters obtained by \citet{bruntt12} and \citet{thygesen12}, and those derived by two other methods has recently been presented by \citet{molenda_zakowicz13}. For the reader interested in the differences arising from the use of different spectroscopic methods, see, e.g., \citet{gillon_magain06} and \citet{creevey12}. The impact of the neglect of non-LTE effects on the parameters inferred from excitation and ionisation balance of iron is discussed by, e.g., \citet{lind12} and \citet{bensby14}. \begin{table} \scriptsize \caption{Typical 1-$\sigma$ uncertainty of the parameter determination for the seismic targets. When available, the second row gives for a given study the uncertainties in case the gravity is fixed to the seismic value (see Sect.~\ref{sect_adopting_seismic_logg}). References: [1] \citet{gazzano10}; [2] \citet{morel14}; [3] \citet{bruntt12}; [4] \citet{thygesen12}.} \label{tab_uncertainties} \begin{tabular}{p{3.2cm}p{1.8cm}p{2.3cm}p{0.8cm}p{0.9cm}p{0.9cm}p{1.0cm}} \hline\noalign{\smallskip} Type of stars & Magnitude range & Type of data & $\sigma_{\rm \, T_{eff}}$ & $\sigma_{\rm \, \log g}$ & $\sigma_{\rm \, [Fe/H]}$ & Reference\\ \noalign{\smallskip}\svhline\noalign{\smallskip} Stars in CoRoT exofields & 12 $<$ $r'$ $<$ 16 & medium resolution$^a$ & 140 & 0.27 & 0.19 & 1\\ Giants in CoRoT seismofields & 6 $<$ $V$ $<$ 9 & high resolution & 85 & 0.20 & 0.10 & 2\\ & & & 60 & 0.07 & 0.08 & 2\\ Dwarfs in {\it Kepler} field & 7 $<$ $V_{\rm \, T}$ $<$ 10.5 & high resolution & 70 & 0.08 & ... & 3\\ & & & 60 & 0.03 & 0.06 & 3\\ Giants in {\it Kepler} field & 7 $<$ $V$ $<$ 12 & high resolution & 80 & 0.20 & 0.15 & 4\\ \noalign{\smallskip}\hline\noalign{\smallskip} \end{tabular} $^a$ Also small wavelength coverage ($\sim$200 \AA). \end{table}
14
3
1403.5103
A full exploitation of the observations provided by the CoRoT and Kepler missions depends on our ability to complement these data with accurate effective temperatures and chemical abundances. We review in this contribution the major efforts that have been undertaken to characterise late-type, seismic targets based on spectra gathered as part of the ground-based, follow-up campaigns. A specific feature of the spectroscopic studies of these stars is that the gravity can be advantageously fixed to the more accurate value derived from the pulsation spectrum. We will describe the impact that such an approach has on the estimation of T <SUB>eff</SUB> and [Fe/H]. The relevance of red-giant seismic targets for studies of internal mixing processes and stellar populations in our Galaxy will also be briefly discussed.
false
[ "chemical abundances", "accurate effective temperatures", "internal mixing processes", "stellar populations", "Galaxy", "part", "studies", "spectra", "red-giant seismic targets", "the pulsation spectrum", "Kepler", "CoRoT", "late-type, seismic targets", "[Fe/H", "the spectroscopic studies", "the ground-based, follow-up campaigns", "the more accurate value", "T <SUB>eff</SUB", "the CoRoT and Kepler missions", "our Galaxy" ]
7.473825
11.855099
81
437583
[ "Cáceres, C.", "Kabath, P.", "Hoyer, S.", "Ivanov, V. D.", "Rojo, P.", "Girard, J. H.", "Miller-Ricci Kempton, E.", "Fortney, J. J.", "Minniti, D." ]
2014A&A...565A...7C
[ "Ground-based transit observations of the super-Earth GJ 1214 b" ]
19
[ "Instituto de Física y Astronomía, Universidad de Valparaíso, Av. Gran Bretaña 1111, 2360102, Valparaíso, Chile ; Millenium Nucleus \"Protoplanetary Disks in ALMA Early Science\", Universidad de Valparaíso, 2360102, Valparaíso, Chile", "European Southern Observatory, Av. Alonso de Córdova 3107, Vitacura, 19001, Santiago, Chile", "Instituto de Astrofísica de Canarias, Vía Láctea s/n, 38200 - La Laguna, Tenerife, Canary Islands, Spain; Department of Astrophysics, University of La Laguna, Vía Láctea s/n, 38200-La Laguna, Tenerife, Canary Islands, Spain", "European Southern Observatory, Av. Alonso de Córdova 3107, Vitacura, 19001, Santiago, Chile", "Universidad de Chile, Camino El Observatorio 1515, Las Condes, Casilla 36-, D Santiago, Chile", "European Southern Observatory, Av. Alonso de Córdova 3107, Vitacura, 19001, Santiago, Chile", "Department of Physics, Grinnell College, Grinnell, IA, 50112, USA", "University of California, Santa Cruz, Department of Astronomy &amp; Astrophysics, 1156 High Street, Santa Cruz, CA, 95064, USA", "Pontificia Universidad Católica de Chile, Departamento de Astronomía y Astrofísica, Av. Vicuña Mackenna 4860, 782-0436 Macul, Santiago, Chile; Specola Vaticana, 00120, Vatican City State, Italy" ]
[ "2015A&A...578A.133A", "2017A&A...600A..10M", "2017A&A...608A.120A", "2017ApJ...834..151R", "2018A&A...612A.105D", "2018AJ....155..208A", "2018ApJ...863..165G", "2018PASA...35...47M", "2018arXiv180502696N", "2018exha.book.....P", "2019ApJ...871...56M", "2019MNRAS.483.3781K", "2020ExA....50....1G", "2022A&A...662A..38N", "2022ApJS..258...40K", "2022ExA....53..547K", "2023ApJS..265....4K", "2024MNRAS.530.2475Y", "2024arXiv240415430K" ]
[ "astronomy" ]
0
[ "planetary systems", "techniques: photometric", "techniques: spectroscopic", "planets and satellites: atmospheres", "stars: individual: GJ 1214", "Astrophysics - Earth and Planetary Astrophysics", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1992StaSc...7..457G", "1992nrfa.book.....P", "1993SPIE.1946..667D", "1998Msngr..91....9M", "1999MNRAS.304..793C", "2002ApJ...580L.171M", "2004PhRvD..69j3501T", "2005AJ....129.1706F", "2005ApJ...631.1198G", "2005MNRAS.364..649F", "2006A&A...459..249G", "2006ApJ...642..505F", "2006ApJ...652.1715H", "2007MNRAS.380.1230C", "2008ApJ...683.1076W", "2009A&A...506..359G", "2009A&A...507..481C", "2009Natur.462..891C", "2010A&A...511A...3G", "2010ApJ...716.1208R", "2010ApJ...716L..74M", "2010ApJ...720L.215S", "2010Natur.468..669B", "2010PASP..122..935E", "2011A&A...529A..75C", "2011A&A...530A...5C", "2011ApJ...730...82C", "2011ApJ...731..123K", "2011ApJ...731L..40D", "2011ApJ...733....2N", "2011ApJ...733...53H", "2011ApJ...736...12B", "2011ApJ...736...78C", "2011ApJ...736..132C", "2011ApJ...743...92B", "2012A&A...538A..46D", "2012A&A...544A..41M", "2012ApJ...745....3M", "2012ApJ...747...35B", "2012ApJ...748...22H", "2012ApJ...756..176H", "2013A&A...549A..10H", "2013A&A...551A..48A", "2013A&A...559A..33C", "2013ApJ...765..127F", "2013ApJ...771..109D", "2013ApJ...773..144N", "2013ApJ...775...33M", "2013ApJ...776...49C", "2013MNRAS.431.1669T", "2013PASJ...65...27N", "2014A&A...563A..21G", "2014MNRAS.438.2395W", "2014Natur.505...69K" ]
[ "10.1051/0004-6361/201321087", "10.48550/arXiv.1403.2723" ]
1403
1403.2723_arXiv.txt
The handful of known transiting extrasolar super-Earths have been mostly found by space missions such as {\it Kepler}, {\it CoRoT}, or {\it MOST}, with the notable exception of GJ\,1214\,b (2.7\,$R_{\oplus}$, 6.55\,$M_{\oplus}$), which was discovered by the ground-based transiting survey {\it MEarth} around a bright nearby M-star \citep{charbonneau_etal2009}. Despite the relatively small radius of GJ\,1214\,b, the small stellar radius of its host star results in a $\sim$1.5\% transit depth, making it one of the super-Earths best suited for follow-up studies. The recent theoretical predictions for the atmosphere of GJ\,1214\,b currently offer the most plausible models that fit the planet's mass, radius, and irradiation level: a rocky or an icy core with a nebular hydrogen-helium envelope, that is, a mini-Neptune, a rocky planet with an outgassed hydrogen atmosphere, or a core with a heavy (up to 45\% of the planet mass) hot water vapor envelope with high molecular mass \citep{rogers_seager_2010, nettelmann_etal2011}, as well as a planet with a cloudy or hazy atmosphere, with a high mean molecular mass composition \citep{morley_etal2013}. GJ\,1214\,b has been the target of an intensive observing campaign. \citet{bean_etal2010} reported a flat optical transmission spectrum that ruled out a cloud-free mini-Neptune model. This conclusion was supported by additional transit observations of \citet{bean_etal2011}, \citet{berta_etal_2012}, and \citet{crossfield_etal_2011} in the near-infrared (NIR) region, and \citet{desert_etal2011} in the mid-infrared. However, \citet{croll_etal2011} measured a deeper 2.2\,$\mu$m transit, implying H$_2$ absorption, consistent with a mini-Neptune model. Transient achromatic haze might reconcile these observations, but the close dates of the $K$-band observations of \citet{croll_etal2011} and \citet{crossfield_etal_2011} make it unlikely, and as \citet{berta_etal_2012} pointed out, there is no known source of achromatic haze. Strangely, \citet{narita_etal2012} reported $J-$ and $H-$band observations consistent with a flat featureless transmission spectrum, but shallower $K_S$-band transmission depth, in disagreement with the result of \citet{croll_etal2011}. \citet{fraine_etal_2012} also reported constant planetary atmosphere radii at $I+z$, 3.6\,$\mu$m, and 4.5\,$\mu$m. Finally, \citet{murgas_etal2012} found a radius of GJ\,1214\,b at the H$\alpha$ line higher than the radii measured at nearby continua, but the difference is not statistically significant. The latest observations reported by \citet{teske_etal2013}, \citet{colon_gaidos2013}, \citet{wilson_etal2013}, \citet{gillon_etal2013}, \citet{demooij_etal2013}, and \citet{kreidberg_etal2014} have agreed in showing a featureless spectrum, favoring high mean molecular mass compositions. In that context, recently \citet{kreidberg_etal2014} ruled out the cloud-free high molecular weight atmosphere scenario for GJ\,1214b with a high statistical significance. Here we report new optical and NIR ground-based observations of GJ\,1214\,b transits, aiming to independently verify these results. The paper is organized as follows: Sect.\,\ref{sec:obs} presents our observations, Sect.\,\ref{sec:analysis} describes the analysis, Sect.\,\ref{sec:disc} presents a discussion of our results, and Sect.\,\ref{sec:summary} is a summary of the main points of this paper.
} \subsection{Spectrophotometric observations} The careful analysis shows that SOFI is fairly stable in terms of differential flux losses, and the duty cycle was extremely high -- 95-97\%, that is, we lost only 3-5\% of the time for detector readout, fits-file merging, and transferring. Unfortunately, the loss of the third night because of poor weather conditions undermined the analysis of our spectroscopic data, and the brightness of the exoplanet host was marginal. We still present the data here as a benchmark and reference for similar future observations, to demonstrate the potential of long-slit spectrographs at 4-m class telescopes to study extrasolar planets \citep[e.g.][]{crossfield_etal2013}. SOFI has a major advantage over other spectrographs: its long slit length of $\sim 5\arcmin$, which facilitates finding a suitable comparison star, which is necessary to control the systematics (see discussion in previous section and Fig.~\ref{fig:rednoisesofi}). Therefore, we conclude that SOFI is suitable for studying exoplanet atmospheres, if the planets orbit brighter stars. \subsection{Photometric GJ\,1214\,b atmosphere} A variety of measurements of the effective planet-to-star radius ratio at various wavelengths have provided conflicting clues on the composition of the GJ\,1214b's atmosphere. The first detections of this planet suggested that its radius is too large to be explained by a solid (pure rock or pure ice) composition \citep[e.g.,][]{charbonneau_etal2009}, implying the presence of a significant gaseous atmosphere. Depending on the assumptions made for the composition of the planet's interior, this atmosphere could be composed primarily of hydrogen, water, or some combination thereof, a hypothesis that cannot be probed from a single radius measurement. More recent detections at different wavelengths have suggested a high molecular weight atmosphere with a probable dominance of water, which would show shallow $K$-band depths \citep[e.g.,][]{bean_etal2010, desert_etal2011, bean_etal2011, berta_etal_2012}. At the same time, \citet{croll_etal2011} reported a deep $K_S$-band transit, which favors a low molecular weight atmosphere, with an hydrogen-rich component; this result was accompanied by the relatively deep $K_S$-band transit reported by \citet{demooij_etal2012}. Meanwhile, \citet{crossfield_etal_2011} reported that the hydrogen-rich atmosphere is not favored by their results. As noted by \citet{howe_burrows2012}, the shorter wavelength measurements reported by \citet{bean_etal2011} and \citet{demooij_etal2012} favor hydrogen-rich atmospheres as well. \citet{nettelmann_etal2011} considered different models for GJ\,1214b's interior, inferring metal-rich H/He atmospheres are the most plausible models, and suggesting an H/He/H$_2$O model with a high water mass fraction for the atmosphere of the planet. \begin{figure*}[!t] \centering{ \includegraphics[width=0.47\textwidth]{fig11a.eps} \hspace{20pt} \includegraphics[width=0.47\textwidth]{fig11b.eps}} \caption{\label{fig:zoomin} {\sl Left:} A zoom-in from Fig.~\ref{fig:models} for the optical region around our $I$-Bessel measurements. {\sl Right}: The $K$-band region of spectra around our $2.14\,\mu$m observation. Our measurement points are represented by dark circles, while gray points follow the description in Fig.~\ref{fig:models}. A color version of this plot can be found in the electronic version of the paper.} \end{figure*} There is no individual theoretical model that accounts for all the observational data, but atmospheres that either have a water-rich composition (more than 70\% by mass) or a thick layer of clouds or hazes in the upper atmosphere are the best-suited interpretations of current data \citep[e.g.,][]{fortney_2005, miller-ricci_fortney2010, nettelmann_etal2011, howe_burrows2012, miller-ricci_etal2012}. The last option has been studied by \citet{morley_etal2013}, who found that in an enhanced metallicity atmosphere, clouds that formed either in chemical equilibrium or nonequilibrium frameworks can reproduce current observations. The authors also pointed out that hydrocarbon haze produced by photochemistry can flatten the GJ\,1214\,b spectrum. Some studies have been performed to try to solve the discrepancies between the models. \citet{murgas_etal2012} used GTC tunable filters to attempt the detection of the H$\alpha$ signature during transits of GJ\,1214b, which yielded a nondetection, consistent with the featureless transmission spectra presented by \citet{bean_etal2011}. Finally, \citet{kreidberg_etal2014} have provided strong evidence of the presence of clouds in the atmosphere of GJ\,1214\,b, based on HST transmission spectra. They reported the significant detection of a featureless spectrum, ruling out the cloud-free high molecular weight atmosphere hypothesis. Figure~\ref{fig:models} shows all current observational data that have a wavelength-dependent radius ratio, where the models correspond to updated best-fit atmosphere models proposed in \citet{miller-ricci_etal2012}. Photometric data were obtained from \citet{demooij_etal2012}, \citet{carter_etal2011}, \citet{bean_etal2011}, \citet{murgas_etal2012}, \citet{croll_etal2011}, \citet{desert_etal2011}, \citet{narita_etal2013,narita_etal2012}, \citet{teske_etal2013}, \citet{colon_gaidos2013}, \citet{wilson_etal2013}, \citet{gillon_etal2013}, and \citet{demooij_etal2013}. Transmission spectroscopy measurements were obtained from \citet{berta_etal_2012}, \citet{bean_etal2011}, and \citet{kreidberg_etal2014}. Our measurement at $0.87 \mu m$ agrees well with current measurements at the shorter wavelengths, also suggesting the presence of the Rayleigh scattering tail argued in \citet{howe_burrows2012}. On the other hand, our NIR detection at $2.14\,\mu m$ shows a moderate depth that disagrees with that of \citet{croll_etal2011} and the $K_S$-band detection of \citet{desert_etal2011}. Recently, \citet{narita_etal2012} reported simultaneous $J$, $H$, and $K_S$-band transit depths for transits of GJ\,1214\,b. Of particular interest is their shallow detection at $2.16\,\mu m$, which strongly disagrees with the deeper measurements. Our $2.14\,\mu m$ detection is consistent with the detections of \citet{narita_etal2012} and \citet{bean_etal2011} and the $K_C$ detection of \citet{desert_etal2011}.% For better orientation, Fig.~\ref{fig:zoomin}, left, is a zoom-in of Fig.~\ref{fig:models} with a focus on the optical region around our {\it SOI} measurement, while the same Fig.~\ref{fig:zoomin}, right, shows the near-IR region around our {\it OSIRIS} measurement. For both panels, relevant measurements by the above mentioned groups are presented for comparison. Finally, we would like to point out a discrepancy between the narrow-band and broad-band photometric measurements at similar wavelengths (see Fig.~\ref{fig:zoomin}). Our measurement at $2.14\,\mu$m contrasts with the broad band measurements by~\citet{croll_etal2011} and \citet{narita_etal2012} and agrees well with the narrow-band measurement of \citet{demooij_etal2012} at $2.27\,\mu$m. \begin{figure}[!pt] \begin{center} \includegraphics[width=0.96\columnwidth]{fig12.eps} \end{center} \caption{\label{fig:timing} {\sl Top:} The observed-minus-calculated (O-C) diagram for the whole set of transit timing available to date in the literature, calculated on the ephemeris given by \citet{bean_etal2011}. Our measurements are drawn as filled squares and in addition highlighted by colors in the online version. The dashed line represents the new ephemeris calculated in this work. {\sl Bottom:} The same data set after removing the new ephemeris, including its 1-$\sigma$ errors (dashed lines).} \end{figure} \subsection{Transit-timing observations} The timing information in the raw images was converted from MJD to BJD (TDB) to determine the final individual transit timing, following the prescriptions in \citet{eastman_etal2010}. Our photometric data show no significant deviations from a constant period, which is consistent with what has been found by \citet{carter_etal2011} and \citet{berta_etal_2012}, and has been recently confirmed by \citet{harpsoe_etal2013}, who used a Bayesian approach to infer that a TTV is unlikely to be present in the GJ\,1214b data. We collected all available transit-timing data from the literature, which we combined with our measurements to refine the ephemeris of the GJ1214b system, with parameters $T_0 = 2454934.917003 \pm 0.000023$\,BJD and $P = 1.580404599 \pm 0.000000056$. Timing data from \citet{demooij_etal2012}, \citet{Kundurthy_etal2010}, \citet{carter_etal2011}, \citet{murgas_etal2012}, \citet{bean_etal2011}, \citet{charbonneau_etal2009}, \citet{berta_etal2011}, \citet{berta_etal_2012}, \citet{croll_etal2011}, \citet{desert_etal2011}, \citet{sada_etal2010}, \citet{narita_etal2012}, \citet{harpsoe_etal2013}, \citet{teske_etal2013}, \citet{gillon_etal2013}, \citet{colon_gaidos2013}, \citet{fraine_etal2013}, and our measurements are shown in Fig.~\ref{fig:timing}. Based on the timing analysis of GJ1214b transits, we put strong constraints on the mass of an additional body in the system, especially in mean motion resonances (MMRs) with the transiting exoplanet. Using dynamical simulations with the MERCURY N-body orbital integrator \citep{chambers_1999}, we determined the mass of an orbital perturber as a function of the distance from the star. To run the simulations we followed the same procedure as described in detail in \citet{hoyer_etal2011}. We used the updated physical parameters of the system reported by \citet{anglada-escude_etal2012} as input for the simulations. For the perturber body, we explored a wide range of masses ($0.5$\,$M_{\oplus}$ - $1000$\,$M_{\oplus}$) and orbital distances ($0.0015$\,AU - $0.055$\,AU), searching for the masses that produce an rms of $\sim30$\,s in the calculated central time of the transits of GJ1214b during the ten years of integration time we used. The results of these dynamical simulations are presented in Fig.~\ref{fig:MvsA} where a region of unstable orbits is marked by the gray strip. For all the other stable orbits we determined the perturber mass that would produce a TTV rms of 30\,s (represented by the solid line). In the 1:2, 2:3 (interior), and in the 3:2, 2:1 (exterior) MMRs (vertical lines in Fig.~\ref{fig:MvsA}) the upper mass limits we obtained correspond to $\sim0.5~M_{\oplus}$. These are better constraints in the mass of a possible companion of GJ\,1214\,b than the limits imposed by radial velocity measurements. \begin{figure}[t] \begin{center} \includegraphics[width=0.96\columnwidth]{fig13.eps} \end{center} \caption{\label{fig:MvsA} Upper mass limits for the system GJ\,1214b based on numerical simulations. The dashed line represents the limits imposed by the radial velocity measurements.} \end{figure} } GJ\,1214\,b is undoubtedly one of the most intriguing objects and the first of its class with an extensively studied atmosphere. The near-IR and optical photometric measurements we present here provide additional evidence for a rather flat featureless spectrum, indicating either a metal-rich, or a cloudy or hazy atmosphere. The TTV analysis of our data combined with previous data sets by other groups have confirmed the constant value of the planetary orbital period. All observations reported here were performed with 4m class telescopes and prove the suitability of such facilities for high-precision photometry. Furthermore, we encourage new spectroscopic measurements especially in the $H$ and $K_{s}$ bands of the spectra. Particularly suitable instruments are either multi-object spectrographs or, as presented in this paper, long-slit spectrographs with very wide and long slits capable of simultaneously monitoring a comparison star. In the latter case, 4m class telescope usage as for GJ1214b may be challenging, but might provide very interesting results for multiple events.
14
3
1403.2723
Context. GJ 1214 b is one of the few known transiting super-Earth-sized exoplanets with a measured mass and radius. It orbits an M-dwarf, only 14.55 pc away, making it a favorable candidate for follow-up studies. However, the composition of GJ 1214 b's mysterious atmosphere has yet to be fully unveiled. <BR /> Aims: Our goal is to distinguish between the various proposed atmospheric models to explain the properties of GJ 1214 b: hydrogen-rich or hydrogen-He mix, or a heavy molecular weight atmosphere with reflecting high clouds, as latest studies have suggested. <BR /> Methods: Wavelength-dependent planetary radii measurements from the transit depths in the optical/NIR are the best tool to investigate the atmosphere of GJ 1214 b. We present here (i) photometric transit observations with a narrow-band filter centered on 2.14 μm and a broad-band I-Bessel filter centered on 0.8665 μm, and (ii) transmission spectroscopy in the H and K atmospheric windows that cover three transits. The photometric and spectrophotometric time series obtained were analyzed with MCMC simulations to measure the planetary radii as a function of wavelength. We determined radii ratios of 0.1173<SUB>-0.0024</SUB><SUP>+0.0022</SUP> for I-Bessel and 0.11735<SUB>-0.00076</SUB><SUP>+0.00072</SUP> at 2.14 μm. <BR /> Results: Our measurements indicate a flat transmission spectrum, in agreement with the last atmospheric models that favor featureless spectra with clouds and high molecular weight compositions. <P />Based on observations obtained at the Southern Astrophysical Research (SOAR) Telescope, which is a joint project of the Ministério da Ciência, Tecnologia, e Inovação (MCTI) da República Federativa do Brasil, the US National Optical Astronomy Observatory (NOAO), the University of North Carolina at Chapel Hill (UNC), and Michigan State University (MSU). SofI results are based on observations made with ESO Telescopes at the La Silla Paranal Observatory under programme ID 087.C-0509.Tables of the lightcurve data are only available at the CDS via anonymous ftp to <A href="http://cdsarc.u-strasbg.fr">http://cdsarc.u-strasbg.fr</A> (ftp://130.79.128.5) or via <A href="http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/565/A7">http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/565/A7</A>
false
[ "high molecular weight compositions", "Michigan State University", "high clouds", "latest studies", "da República Federativa", "anonymous ftp", "wavelength", "Chapel Hill", "North Carolina", "radius", "GJ", "programme ID 087.C-0509.Tables", "clouds", "a heavy molecular weight atmosphere", "República Federativa", "MSU", "the US National Optical Astronomy Observatory", "UNC", "featureless spectra", "NOAO" ]
7.114124
14.277979
108
483169
[ "Gutiérrez, Claudia P.", "Anderson, Joseph P.", "Hamuy, Mario", "González-Gaitán, Santiago", "Folatelli, Gastón", "Morrell, Nidia I.", "Stritzinger, Maximilian D.", "Phillips, Mark M.", "McCarthy, Patrick", "Suntzeff, Nicholas B.", "Thomas-Osip, Joanna" ]
2014ApJ...786L..15G
[ "H<SUB>α</SUB> Spectral Diversity of Type II Supernovae: Correlations with Photometric Properties" ]
69
[ "Millennium Institute of Astrophysics, Casilla 36-D, Santiago, Chile; Departamento de Astronomía, Universidad de Chile, Casilla 36-D, Santiago, Chile;", "Departamento de Astronomía, Universidad de Chile, Casilla 36-D, Santiago, Chile; European Southern Observatory, Alonso de Córdova 3107, Casilla 19, Santiago, Chile", "Millennium Institute of Astrophysics, Casilla 36-D, Santiago, Chile; Departamento de Astronomía, Universidad de Chile, Casilla 36-D, Santiago, Chile", "Millennium Institute of Astrophysics, Casilla 36-D, Santiago, Chile; Departamento de Astronomía, Universidad de Chile, Casilla 36-D, Santiago, Chile", "Institute for the Physics and Mathematics of the Universe (IPMU), University of Tokyo, 5-1-5 Kashiwanoha, Kashiwa, Chiba 277-8583, Japan", "Carnegie Observatories, Las Campanas Observatory, Casilla 601, La Serena, Chile", "Department of Physics and Astronomy, Aarhus University, Ny Munkegade 120, DK-8000 Aarhus C, Denmark", "Carnegie Observatories, Las Campanas Observatory, Casilla 601, La Serena, Chile", "Observatories of the Carnegie Institution for Science, Pasadena, CA 91101, USA", "George P. and Cynthia Woods Mitchell Institute for Fundamental Physics and Astronomy, Department of Physics and Astronomy, Texas A&amp;M University, College Station, TX 77843, USA", "Carnegie Observatories, Las Campanas Observatory, Casilla 601, La Serena, Chile" ]
[ "2014ApJ...786...67A", "2014MNRAS.441..671A", "2014MNRAS.445..554F", "2015A&A...582A...3G", "2015AJ....150..156R", "2015ApJ...804...28G", "2015ApJ...806..160B", "2015ApJ...815..121D", "2015MNRAS.448.2608V", "2015MNRAS.449.1753P", "2015MNRAS.449.1876S", "2015MNRAS.451.2212G", "2015PASA...32...19A", "2016A&A...589A.110A", "2016AJ....151...33G", "2016AJ....151...47R", "2016ApJ...820...33R", "2016ApJ...823..127N", "2016MNRAS.461.2003Y", "2016MNRAS.462..137T", "2017ApJ...838...28M", "2017ApJ...841..127M", "2017ApJ...850...89G", "2017ApJ...850...90G", "2017ApJ...851..138D", "2017hsn..book..239A", "2017suex.book.....B", "2018A&A...611A..25G", "2018ApJ...853...62T", "2018ApJ...862..107B", "2018MNRAS.476.4592D", "2018MNRAS.478.3776D", "2018MNRAS.480.2475S", "2019A&A...629A..57M", "2019A&A...631A...8H", "2019ApJ...887....4D", "2019MNRAS.485.5120B", "2019MNRAS.486.2850D", "2019MNRAS.488.4239P", "2019MNRAS.490.2799D", "2019NatAs...3..717M", "2020ApJ...890..177K", "2020ApJ...895...31B", "2020MNRAS.493.1761R", "2020MNRAS.499..974G", "2021ApJ...909..145D", "2021ApJ...913...55H", "2021MNRAS.504.1009D", "2021MNRAS.505.3664N", "2021MNRAS.505.4890L", "2021SSRv..217...54Z", "2021arXiv210604375L", "2022A&A...660A..40M", "2022A&A...660A..41M", "2022A&A...660A..42M", "2022A&A...660L...9D", "2022ApJ...930...34T", "2022NatAs...6.1452A", "2023A&C....4400715D", "2023ApJ...942...17M", "2023MNRAS.523.5315P", "2023Univ....9..201L", "2024A&A...685A.169D", "2024ApJ...961..247S", "2024MNRAS.528.3092L", "2024arXiv240104027M", "2024arXiv240500113T", "2024arXiv240512504H", "2024arXiv240520989S" ]
[ "astronomy" ]
16
[ "supernovae: general", "Astrophysics - High Energy Astrophysical Phenomena", "Astrophysics - Cosmology and Nongalactic Astrophysics", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1979A&A....72..287B", "1981ApJ...244..780B", "1983Ap&SS..89...89L", "1985SvAL...11..145L", "1993ApJ...414..712P", "1994A&A...282..731P", "1995MNRAS.277..106C", "1996AJ....111.1660S", "2000ApJ...545..444B", "2001ApJ...553..861L", "2002AJ....124.2490L", "2002ApJ...566L..63H", "2002ApJ...572..932P", "2002PASP..114...35L", "2003ApJ...582..905H", "2004MNRAS.347...74P", "2005coex.conf..195P", "2006MNRAS.370.1752P", "2006MNRAS.372.1315S", "2006PASP..118....2H", "2007ApJ...662.1136C", "2009ApJ...692..839M", "2009ApJ...704..306K", "2010ApJ...715..541A", "2011MNRAS.417..261I", "2012ApJ...756L..30A", "2012MNRAS.422.1122I", "2013A&A...555A.142I", "2013ApJ...773...53F", "2014ApJ...786...67A" ]
[ "10.1088/2041-8205/786/2/L15", "10.48550/arXiv.1403.7089" ]
1403
1403.7089_arXiv.txt
Type II Supernovae (SNe II) are produced by the final explosion of massive ($>8$ M$_\odot$) stars. They retain a significant part of their hydrogen envelope at the time of the explosion, and hence their spectra show strong Balmer lines. Studies of the variety of SNe~II have relied on photometric analysis, cataloging this group in two sub-classes according to the shape of the light curve: SNe with a plateau (quasi-constant luminosity for a period of a few months) are classified as SNe IIP, while SNe with steeper declining linear light curves as SNe IIL \citep{Barbon79}. However, despite the role played by SNe~II in stellar evolution, the impact on their environments and their importance as standardized candles, an overall picture describing the physics which underpins their diversity is lacking. It has been suggested that SNe~IIL are produced by progenitors which explode with smaller mass H envelopes, which then lead to SNe with more linearly declining light curves and shorter or non-existent `plateaus' \citep{Popov93}. Indeed, this was argued to be the case for the prototype SN~IIL 1979C \citep{Branch81}. This would imply that SNe IIL progenitors suffer from a higher level of mass-loss than their IIP counterparts. In addition, a number of SNe~IIL have shown evidence for circumstellar (CSM) interaction at late times (e.g. SN 1986E, \citealt{Cappellaro95a}; SN~1979C, \citealt{Milisavljevic09}), which has been interpreted as evidence of interaction of the ejecta with the pre-supernova CSM (see e.g. \citealt{Sahu06}, \citealt{Inserra13}). However, a number of authors have also claimed evidence for signs of CSM interaction in SNe~IIP (e.g. SN~1999em, \citealt{Pooley02}; SN~2004et, \citealt{Kotak09}; SN~2007od, \citealt{Inserra11}, \citealt{Andrews10}; SN~2009bw, \citealt{Inserra12b}). \\ \indent In recent years many individual studies have been published focusing on particular properties of individual SNe, but few statistical studies where the spectral and photometric properties have been directly related are available. \citet{Patat94} found correlations and anti-correlations between the maximum B-band magnitude ($M_{max}^B$), the color at maximum ($(B-V)_{max}$) and the ratio of absorption to emission ($e/a$) in H$_{\alpha}$, concluding that SNe IIL have shallower P-Cygni profiles (larger $e/a$ values) than SNe IIP. \citet{Hamuy02L} analysed 17 SNe IIP and found that SNe with brighter plateaus have higher expansion velocities. Similar results were found by \citet{Pastorello04} with four SNe II, who concluded that low lumino\-sity SNe have narrow spectral lines indicating low expansion velocities. \citet{Hamuy03} used observations together with the analytical models of \citet{Litvinova83, Litvinova85} to derive physical SN~IIP properties. He found that more massive progenitors produce more energetic explosions and in turn produce more nickel. These results were confirmed by \citet{Pastorello03} with a heterogeneous group of SNe~II that share a very wide range of physical properties.\\ \indent Despite the above results, it is currently unclear whether underlying spectral and photometric relations exist for the whole ensemble of SN~II events. Therefore, here we attempt to remedy this situation by presenting an initial statistical analysis of various spectroscopic and photometric properties of a large sample of SNe~II.\\ \indent In this letter we present results showing the diversity of H$_{\alpha}$ P-Cygni profiles, and relations between spectral and photometric parameters for a sample of 52 SNe. The letter is organized as follows. In \S\ 2 we outline our SN sample and we define the measurements, then in \S\ 3 we present the results. In \S\ 4 possible physical explanations of those results are discussed, and finally in \S\ 5 we list our conclusions. We note that a detailed analysis of the $V$-band light curve properties of the currently analyzed sample of SN~II is being presented in Anderson et al. (submitted, hereafter A14).
\begin{itemize} \item $a/e$ is an important parameter describing the spectral diversity of SNe~II. \item SNe with low $a/e$ values appear to have high $H_{\alpha}$ velocities and decline rates, are brighter and have a smaller \textit{OPTd} values. \item While any definitive spectral distinction between IIP and IIL is not clear, SNe with higher \textit{$s_{2}$} values (i.e. more `linear' SNe) have smaller $a/e$ values, have higher H$_{\alpha}$ velocities, and are more luminous. \item We speculate that the envelope mass retained before explosion and the density gradient play a very important role to determine the differences of H$_{\alpha}$ P-Cygni profile. \item CSM interaction could also be a cause of the change in the P-Cygni profiles, suggesting that faster declining SNe have more intense interactions. \end{itemize} This paper presented a first analysis of SN~II spectral from CSP. The full analysis of that sample (optical and near IR photometry and spectroscopy) promises to significantly further our knowledge of the SN~II phenomenon.
14
3
1403.7089
We present a spectroscopic analysis of the H<SUB>α</SUB> profiles of hydrogen-rich Type II supernovae. A total of 52 Type II supernovae having well-sampled optical light curves and spectral sequences were analyzed. Concentrating on the H<SUB>α</SUB> P-Cygni profile we measure its velocity from the FWHM of the emission and the ratio of absorption to emission (a/e) at a common epoch at the start of the recombination phase, and search for correlations between these spectral parameters and photometric properties of the V-band light curves. Testing the strength of various correlations we find that a/e appears to be the dominant spectral parameter in terms of describing the diversity in our measured supernova properties. It is found that supernovae with smaller a/e have higher H<SUB>α</SUB> velocities, more rapidly declining light curves from maximum during the plateau and radioactive tail phase, are brighter at maximum light, and have shorter optically thick phase durations. We discuss possible explanations of these results in terms of physical properties of Type II supernovae, speculating that the most likely parameters that influence the morphologies of H<SUB>α</SUB> profiles are the mass and density profile of the hydrogen envelope, together with additional emission components due to circumstellar interaction. <P />This paper includes data gathered with the 6.5 m Magellan Telescopes located at Las Campanas Observatory, Chile; and the Gemini Observatory, Cerro Pachon, Chile (Gemini Program GS-2008B-Q-56). Based on observations collected at the European Organisation for Astronomical Research in the Southern Hemisphere, Chile (ESO Programmes 076.A-0156, 078.D-0048, 080.A-0516, and 082.A-0526).
false
[ "radioactive tail phase", "maximum light", "Type II supernovae", "spectral sequences", "photometric properties", "physical properties", "additional emission components", "Las Campanas Observatory", "circumstellar interaction", "Chile", "Gemini Program", "emission", "supernovae", "α</SUB", "082.A-0526", "shorter optically thick phase durations", "hydrogen-rich Type II supernovae", "Q-56", "> velocities", "Cerro Pachon" ]
2.815278
4.670362
-1
745751
[ "Ogrean, G. A.", "Brüggen, M.", "van Weeren, R.", "Röttgering, H.", "Simionescu, A.", "Hoeft, M.", "Croston, J. H." ]
2014MNRAS.440.3416O
[ "Multiple density discontinuities in the merging galaxy cluster CIZA J2242.8+5301" ]
53
[ "Hamburger Sternwarte, University of Hamburg, Gojenbergsweg 112, D-21029 Hamburg, Germany", "Hamburger Sternwarte, University of Hamburg, Gojenbergsweg 112, D-21029 Hamburg, Germany", "Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138, USA", "Leiden Observatory, Leiden University, P.O. Box 9513, NL-2300 RA Leiden, the Netherlands", "Japan Aerospace Exploration Agency, 3-1-1 Yoshinodai, Sagamihara, Kanagawa 229-8510, Japan", "Thüringer Landessternwarte, Sternwarte 5, D-07778 Tautenburg, Germany", "School of Physics and Astronomy, University of Southampton, Southampton SO17 1BJ, UK" ]
[ "2014MNRAS.443.2463O", "2014MNRAS.445.1213S", "2015A&A...582A..87A", "2015ApJ...802...46J", "2015ApJ...805..143D", "2015ApJ...809..186K", "2015ApJ...812...49H", "2015MNRAS.450..630S", "2015MNRAS.450..646S", "2015MNRAS.452.2731S", "2015PASJ...67..114O", "2016ApJ...818..204V", "2016ApJ...823...13K", "2016FrASS...2....7M", "2016JKAS...49..145K", "2016MNRAS.458..681D", "2016MNRAS.462.2014D", "2016PASJ...68S..23I", "2016arXiv160607433S", "2017A&A...600A..18K", "2017ApJ...840...42K", "2017ApJ...841...46M", "2017Galax...5...16V", "2017JKAS...50...93K", "2017MNRAS.465.2916S", "2017MNRAS.469.1476S", "2017MNRAS.470.4638R", "2017MNRAS.471.1107H", "2017MNRAS.471.4587D", "2017MNRAS.472.3605L", "2018A&A...619A..68T", "2018ApJ...857...26H", "2018ApJ...865...24D", "2018MNRAS.478.4922Z", "2018MNRAS.479..553S", "2019A&A...629A..31U", "2019ApJ...874...71P", "2019SSRv..215...16V", "2020A&A...634A..64B", "2020ApJ...903..103W", "2020ApJ...905L..22S", "2020MNRAS.493.2306B", "2020MNRAS.495.4392N", "2020MNRAS.495.5014B", "2020MNRAS.498.1628L", "2021ApJ...911....3D", "2021ApJ...914...73Z", "2021Galax...9...91Z", "2022hxga.book...13W", "2023A&ARv..31....4R", "2023MNRAS.519..548B", "2023arXiv230712423H", "2024MNRAS.527.1194K" ]
[ "astronomy" ]
7
[ "shock waves", "galaxies: clusters: individual: CIZA J2242.8+5301", "X-rays: galaxies: clusters", "Astrophysics - Astrophysics of Galaxies", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
[ "1979ApJ...228..939C", "1983RPPh...46..973D", "1989MNRAS.239..995K", "1993MNRAS.261..921K", "1993MNRAS.262..627L", "1994RPPh...57..325K", "1996ApJ...462..563N", "2005ApJ...627..733M", "2007ApJ...662..224K", "2007ApJ...669..729K", "2007MNRAS.375...77H", "2007MNRAS.380L..67B", "2009A&A...504...33V", "2010ApJ...715.1143F", "2010MNRAS.406.1721R", "2010MNRAS.408..199B", "2010Sci...330..347V", "2011A&A...526A..79E", "2011ApJ...728...82M", "2012A&ARv..20...54F", "2012ApJ...756...97K", "2012MNRAS.421.1868V", "2012MNRAS.425L..76B", "2013A&A...556A..44R", "2013AN....334..342O", "2013ApJ...764...82B", "2013ApJ...765...21S", "2013MNRAS.429.2617O", "2013MNRAS.433..812O", "2013PASJ...65...16A" ]
[ "10.1093/mnras/stu537", "10.48550/arXiv.1403.5273" ]
1403
1403.5273_arXiv.txt
\label{s:intro} Galaxy clusters grow via mergers with less massive structures \citep[e.g.,][]{Kauffmann1993,Lacey1993,nfw1996}. When clusters merge, shock fronts are triggered into the intracluster medium (ICM). As shocks propagate outwards, they (re-)accelerate particles to relativistic energies. The high-energy particles interact with intracluster magnetic fields \citep[e.g.,][]{Kronberg1994} to create diffuse sources of synchrotron emission visible at radio frequencies \citep[e.g.,][and references therein]{Feretti2012}. These sources are known as radio relics, and are commonly observed at the cluster periphery \citep[e.g.,][]{Vazza2012}. X-ray observations of clusters hosting radio relics reveal the position, extent, and strength of shock fronts. These results can be compared to radio observations to examine the correspondence between radio relics and shock fronts. Most often, relics roughly trace shock fronts present in the ICM \citep[e.g.,][]{Finoguenov2010,Macario2011,Akamatsu2013,Bourdin2013}. Sometimes, however, shock fronts are found in clusters with no radio relics \citep{Russell2010}, or shocks are significantly offset from the relics \citep{Ogrean2013d}. For a theoretical perspective on these challenges and on other open questions in our understanding of cosmic rays in galaxy clusters, see, e.g., \citet{Brunetti2014}. From an observational standpoint, to understand the connection between radio relics and shock fronts, it is essential to investigate whether the few cases that deviate from the norm are mere exceptions, or more common occurrences in merging clusters. Here, we present results from a 200-ks \chandra\ observation of the merging galaxy cluster CIZA J2242.8+5301. The cluster is located at $z=0.1921$, and has a $0.1-2.4$ keV luminosity of $6.8\times 10^{44}$ erg~s$^{-1}$ \citep{Kocevski2007}. \citet{vanWeeren2010} presented a radio analysis of the cluster, which revealed the presence of a double-relic system and a faint, strongly-elongated halo. Both the relics and the halo are oriented along the N-S direction, which suggests that we are observing two clusters with a N-S-oriented merger axis. Of the three main radio structures, the northern relic stands out due to its spectacular morphology: a length of 2 Mpc vs. a width of only 55 kpc. Due to this extreme shape, the relic was nicknamed the ``Sausage''. The spectral index at the outer edge of the relic predicts a Mach number of $4.6\pm 1.1$ under the assumptions of diffusive shock acceleration (DSA) of thermal particles in the linear test-particle regime. \citet{Akamatsu2013} analysed \suzaku\ observations of the cluster, and measured a temperature drop by a factor of $\sim 3$ across the ``Sausage'', which corresponds to a Mach number of $3.15\pm 0.52_{-1.20}^{+0.40}$ (the first error is the statistical error, while the following two are systematic errors). \xmm\ observations of the cluster revealed a surface brightness discontinuity east of the southern relic, and a merger geometry more complex than a binary head-on collision \citep{Ogrean2013b}. In the following, we further explore the surface brightness morphology and the temperature distribution within the ICM using our \chandra\ observations (PI: Ogrean) and archival \suzaku\ observations (PI: Kawahara). The paper is organised as follows: Section \ref{s:obs} presents the \chandra\ observations and the data reduction. In Section \ref{s:analysis} we analyse the X-ray morphology observed with \chandra, and in Section \ref{s:spectroscopy} we use the \suzaku\ datasets to measure the temperature distribution around the structures identified in the previous section. The results are discussed in Section \ref{s:discussion}. We assume a flat $\Lambda$CDM universe with $H_0=70$ km\,s$^{-1}$\,Mpc$^{-1}$, $\Omega_{\rm M}=0.3$, and $\Omega_{\rm \Lambda}=0.7$. At the redshift of the cluster, 1 arcmin corresponds to 192 kpc. Throughout the paper, quoted errors are $1\sigma$ statistical errors. \begin{figure*} \includegraphics[width=\textwidth]{sxmap-radio.eps} \caption{Surface brightness map in the energy band $0.5-7$ keV. The map was binned by a factor of 4, exposure-corrected, vignetting-corrected, instrumental background-subtracted, and smoothed with a Gaussian of kernel size 3 pixels ($\approx 6$ arcsec). Point sources were removed. Overlaid are 1.4~GHz WSRT radio contours, drawn at $[1,4,16,...]\times 70$ $\mu$Jy/beam.} \label{fig:sxmap-radio} \end{figure*}
\label{s:discussion} Analysing the surface brightness profiles along directions on- and off-relic, we identified density discontinuities in all directions from the cluster centre. These discontinuities correspond to very weak Mach numbers of $\sim 1.3$. The only clear temperature jump is detected across the northern relic; it corresponds to a Mach number of $2.54_{-0.43}^{+0.64}$. However, the temperatures throughout the cluster are very high, larger than 5 keV out to distances as large as 1.5 Mpc from the cluster centre \citep{Ogrean2013b}. These high temperatures are difficult to constrain given the useful energy range of \suzaku\ (as well as \xmm\ and \chandra), roughly $0.5-7$ keV. The large statistical uncertainties on the measured temperatures make it impossible to confirm in temperature the Mach numbers calculated from the surface brightness discontinuities. Yet, small temperature jumps associated with Mach numbers of $\sim 1.3$ cannot be excluded. The density discontinuities identified in the \chandra\ observations of CIZA J2242.8+5301 raise several questions related to particle acceleration at cluster merger shocks: \begin{enumerate} \item How does the temperature jump detected across the northern relic compare with the radio-predicted Mach number? \item What is the nature of the inner discontinuities? \item Why is there no radio emission associated with the detected density edges? \item What is the merger scenario that triggered multiple sequential discontinuities in the ICM? \end{enumerate} We address each of these questions below. \subsection{The shock at the N relic} The outer northern discontinuity in surface brightness is only weakly detected. A density jump was expected here, as it would trace the edge of the ``Sausage'' relic and a temperature jump has already been detected near its location by \citet{Akamatsu2013}. We also detect a temperature jump. However, while the Mach number inferred from the temperature jump is relatively small, $\mathcal{M} = 2.54_{-0.43}^{+0.64}$, the Mach number predicted based on the radio spectral index under the assumptions of diffusive shock acceleration in the test-particle regime \citep{Drury1983} is much larger, $\mathcal{M}=4.6\pm 1.1$ \citep{vanWeeren2010}. The same discrepancy between the radio-predicted and X-ray-derived Mach numbers was observed at the northern relic in 1RXS J0603.3+4214 \citep{Ogrean2013d}. Numerical simulations are required to clarify the reasons for the discrepancy. One possibility is that the Mach number varies across the shock front \citep{Skillman2013}, and the synchrotron emission is more sensitive to high Mach number shocks \citep{HoeftBrueggen2007}. Alternatively, the shock does not only accelerate particles from the thermal pool, but re-accelerates a pre-existing cosmic ray particle population. This scenario has been simulated by \citet{Kang2012}, who showed that the northern relic can be reproduced either by a shock of Mach number $\sim 4$, or by a weaker shock of Mach number $\sim 2$. Other possible explanations are projection effects, underestimation of the postshock temperature due to averaging in a wide partial annulus (width $\sim 450$ kpc), or oblique, rather than parallel, shocks \citep{Kirk1989}. \subsection{The nature of the inner density discontinuities} The nature of the inner density discontinuities is difficult to explain in the absence of temperature jumps. On one hand, it seems rather unlikely that the northern discontinuity is a cold front, given its large distance from the centre. At $\approx 1.5$ Mpc from the centre, a cold front in the northern region of CIZA J2242.8+5301 would be the most distant cold front ever detected. At the moment, the most known distant cold front was found in the cluster Abell 2142, at 1 Mpc from the centre \citep{Rossetti2013}. Moreover, while the \xmm\ temperature in the direction of the northern relic revealed a hint of a temperature increase on the outer side of the discontinuity \citep{Ogrean2013a}, this result is not confirmed by \emph{Suzaku}. The temperature on both sides of the northern discontinuity is $\sim 8-9$ keV, which would also mean that if this discontinuity was associated with a cold front, it would be the hottest of all known cold fronts. Nevertheless, while there are indications that at least the northern discontinuity is more likely a shock front, only more precise temperature measurements can confirm this supposition as well as identify the nature of the other inner density jumps. \subsection{No radio emission at the inner density discontinuities} Interestingly, no diffuse radio emission is detected at either of the inner surface brightness discontinuities. If these discontinuities are associated with shock fronts, then the situation is similar to that in Abell 2146 \citep{Russell2010}. However, unlike in Abell 2146, the shocks are not present only along the merger axis, but also west and east of it. More importantly from a particle acceleration viewpoint, while the shock front in Abell 2146 has a Mach number of $\approx 2$, the inner density discontinuities in CIZA J2242.8+5301 appear to be associated with much weaker shocks. Indeed, DSA of thermal particles at very weak shocks does not efficiently accelerate electrons to cosmic ray energies \citep{Kang2007}. Moreover, for a $\mathcal{M}=1.3$ shock, DSA of thermal particles in the test-particle regime would imply an extremely steep spectral index of $\approx 4$, making a relic, if present, too faint in the radio band. Another possibility is that radio relics require a pre-existing CR particle population \citep[e.g.,][]{Kang2007,Kang2012}. In that case, the lack of radio emission at the inner, W, and E shocks suggests that no pre-existing CR population existed at the shock locations previous to the shock passage. However, even if a pre-existing CR population was present, re-acceleration would boost the electron spectrum only by $3C/[C+2-\delta (C-1)] \approx 2$ for a shock of Mach number $1.3$ and a slope of the pre-existing particle population spectrum $\delta=2-3$ \citep{Markevitch2005}. If the discontinuities are associated with cold fronts rather than with shocks, then radio emission is naturally not expected. \subsection{Sequential shocks in the ICM} If the inner discontinuities are shock fronts, then it is for the first time that multiple sequential merger shocks are observed in a galaxy cluster. In numerical simulations, multiple consecutive shocks have been seen in the accretion region of clusters \citep[e.g.,][]{Vazza2009} and close to an active AGN \citep[][]{Brueggen2007}, but not at moderate distances ($\sim 1$ Mpc) from the cluster centre. However, secondary shocks in a binary cluster merger could stem (i) from a second core passage of the DM cores, or (ii) from the violent relaxation of the newly merged halo. In the former case, the discontinuities are strongest along the merger axis and in the latter case, the shocks can be radial with a weaker variation with angle from the merger axis. In order to better understand how multiple shocks can be produced, we have performed very simple dark matter (DM) + hydrodynamic simulations of binary cluster mergers using the adaptive-mesh refinement code {\sc flash}. For more details on the initial conditions and the physics employed in these simulations see \citet{Brueggen2012}. In these idealised simulations of nearly equal mass mergers, we set two clusters in virial and hydrostatic equilibrium on a collision course with a small impact parameter to destroy axial symmetry. Note that these simple simulations were set up to study the interaction between the gas and the DM and were not tuned to reproduce conditions in CIZA J2242.8+5301. In Fig. \ref{fig:sim} we show the logarithm of the gas density in a cut through the plane that contains the two cluster centres. The outermost shock is produced by the impact of the two gas spheres, which lead to a region of high compression and an ellipsoidal shock waves that run outwards. Meanwhile, the DM halos get compressed and tidal tails develop at the far sides of the DM halos. When the DM cores separate again after core passage, they go through violent relaxation as the total gravitational potential varies rapidly and the DM cores expand swiftly (on the violent relaxation time scale which is the timescale on which the gravitational potential changes). This rapid expansion can cause another set of shocks (usually weaker) through the gas, and this is seen as the inner discontinuity. Curiously, in the literature on cluster mergers, this is almost never mentioned. Violent relaxation is mentioned but not its impact on the ICM, presumable because it is dynamically not important and it has never been observed before. The possibility of two shocks is mentioned, however, in \citet{Birnboim2010}, who studied merging shock fronts in the context of cold fronts in galaxy clusters. \begin{figure} \includegraphics[width=\columnwidth]{denscut.eps} \caption{Logarithm of the gas density in a cut through the plane that contains the two cluster centres. The location of the two sets of discontinuities are indicated by the dashed-line regions.} \label{fig:sim} \end{figure} Finally, the presence of multiple putative shocks could also suggest a more complex merger geometry than a simple binary, head-on collision, and complex subsequent processes within the ICM. We leave the investigation of the origin of these possible shock structures to future papers.
14
3
1403.5273
CIZA J2242.8+5301, a merging galaxy cluster at z = 0.19, hosts a double-relic system and a faint radio halo. Radio observations at frequencies ranging from a few MHz to several GHz have shown that the radio spectral index at the outer edge of the northern relic corresponds to a shock of Mach number 4.6_{-0.9}^{+1.3}, under the assumptions of diffusive shock acceleration of thermal particles in the test-particle regime. Here, we present results from new Chandra observations of the cluster. The Chandra surface brightness profile across the northern relic only hints to a surface brightness discontinuity (&lt;2σ detection). Nevertheless, our reanalysis of archival Suzaku data indicates a temperature discontinuity across the relic that is consistent with a Mach number of 2.54_{-0.43}^{+0.64}, in agreement with previously published results. This confirms that the Mach number at the shock traced by the northern relic is much weaker than predicted from the radio. Puzzlingly, in the Chandra data we also identify additional inner small density discontinuities both on and off the merger axis. Temperature measurements on both sides of the discontinuities do not allow us to undoubtedly determine their nature, although a shock front interpretation seems more likely. We speculate that if the inner density discontinuities are indeed shock fronts, then they are the consequence of violent relaxation of the dark matter cores of the clusters involved in the merger.
false
[ "Radio observations", "shock fronts", "diffusive shock acceleration", "Mach number", "additional inner small density discontinuities", "thermal particles", "new Chandra observations", "several GHz", "the northern relic corresponds", "a faint radio halo", "the northern relic", "Mach", "results", "violent relaxation", "GHz", "a shock front interpretation", "the merger axis", "a surface brightness discontinuity", "frequencies", "archival Suzaku data" ]
14.186375
4.710546
-1
800304
[ "Jenkins, Elizabeth E.", "Manohar, Aneesh V.", "Waalewijn, Wouter J.", "Yadav, Amit P. S." ]
2014PhLB..736....6J
[ "Gravitational lensing of the CMB: A Feynman diagram approach" ]
3
[ "Department of Physics, University of California at San Diego, La Jolla, CA 92093, USA", "Department of Physics, University of California at San Diego, La Jolla, CA 92093, USA", "Nikhef, Theory Group, Science Park 105, 1098 XG, Amsterdam, The Netherlands", "Department of Physics, University of California at San Diego, La Jolla, CA 92093, USA" ]
[ "2014JCAP...09..024J", "2023JCAP...04..041K", "2024arXiv240204309S" ]
[ "astronomy", "physics" ]
5
[ "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
[ "1995PhRvD..52.4307K", "1996ApJ...463....1S", "1996ApJ...469....1K", "1997A&A...322....1B", "1997PhRvD..55.7368K", "1997PhRvL..78.2054S", "1997PhRvL..78.2058K", "1998PhRvD..58b3003Z", "1998PhRvL..81.3067C", "1999PhRvD..60d3504S", "2000PhRvD..62f3510Z", "2002ApJ...574...19C", "2002ApJ...574..566H", "2003NewA....8..231C", "2003PhRvD..67h3002O", "2003PhRvD..67l3507K", "2003PhRvD..68h3002H", "2005CoPhC.168...78H", "2005PhRvD..71d3006K", "2005PhRvD..71l3527C", "2009AIPC.1141...10B", "2009PhRvD..79d3003D", "2009PhRvD..79l3009Y", "2009PhRvD..80b3510G", "2009PhRvL.102k1302K", "2009PhRvL.103e1302P", "2010PhRvD..81f3512Y", "2011JCAP...03..018L", "2011PhRvD..83d3005H", "2012JCAP...01..007N", "2012PhRvD..86l3009Y", "2013PhRvD..88h3517A", "2014JCAP...09..024J", "2014PhRvL.112x1101B" ]
[ "10.1016/j.physletb.2014.07.002", "10.48550/arXiv.1403.2386" ]
1403
1403.2386_arXiv.txt
14
3
1403.2386
We develop a Feynman diagram approach to calculating correlations of the Cosmic Microwave Background (CMB) in the presence of distortions. As one application, we focus on CMB distortions due to gravitational lensing by Large Scale Structure (LSS). We study the Hu-Okamoto quadratic estimator for extracting lensing from the CMB and derive the noise of the estimator up to O (ϕ<SUP>4</SUP>) in the lensing potential ϕ. By identifying the diagrams responsible for the previously noted large O (ϕ<SUP>4</SUP>) term, we conclude that the lensing expansion does not break down. The convergence can be significantly improved by a reorganization of the ϕ expansion. Our approach makes it simple to obtain expressions for quadratic estimators based on any CMB channel, including many previously unexplored cases. We briefly discuss other applications to cosmology of this diagrammatic approach, such as distortions of the CMB due to patchy reionization, or due to Faraday rotation from primordial axion fields.
false
[ "CMB distortions", "primordial axion fields", "Large Scale Structure", "CMB", "quadratic estimators", "distortions", "LSS", "ϕ<SUP>4</SUP", "patchy reionization", "the lensing potential ϕ.", "the lensing expansion", "rotation", "many previously unexplored cases", "Faraday", "any CMB channel", "correlations", "other applications", "a Feynman diagram approach", "expressions", "the Cosmic Microwave Background" ]
12.655147
2.133817
103
813466
[ "Giacinti, G.", "Kachelrieß, M.", "Semikoz, D. V." ]
2014PhRvD..90d1302G
[ "Explaining the spectra of cosmic ray groups above the knee by escape from the Galaxy" ]
38
[ "University of Oxford, Clarendon Laboratory, Oxford OX1 3PU, United Kingdom", "Institutt for fysikk, NTNU, Trondheim N-7491, Norway", "AstroParticle and Cosmology (APC), Paris F-75205, France" ]
[ "2014MNRAS.443..474P", "2014NuPhS.256..197P", "2014PhRvD..89j3002N", "2014PhRvD..90h3002K", "2014PhRvD..90l3010A", "2015ApJ...809L..23S", "2015JETP..120..541K", "2015JPhCS.632a2037K", "2015JPhCS.632a2094G", "2015PhRvD..91h3009G", "2015PhRvD..92h3016G", "2015PhRvL.115r1103K", "2016A&A...591A..71M", "2016APh....72...32N", "2016MNRAS.460.3479T", "2016RAA....16....6B", "2017JCAP...06..046M", "2017PhRvD..96h3006K", "2018JCAP...07..051G", "2018PhRvD..97f3011K", "2019EPJWC.21004003K", "2019FrASS...6...23B", "2019JCAP...01..046B", "2019JPhCS1181a2039K", "2019JPhCS1263a2009S", "2019PrPNP.10903710K", "2020APh...114...22M", "2020PhRvD.101l3023B", "2020PhRvD.102d3025N", "2021PhRvD.104b3017N", "2021PhRvD.104d3010K", "2022AdSpR..70.2696T", "2022JHEAp..35....1V", "2022PhRvD.106f3004B", "2022PhRvD.106l3029G", "2022arXiv221105881K", "2023PhRvD.108j3008N", "2023arXiv230510251G" ]
[ "astronomy", "physics" ]
6
[ "98.70.Sa", "96.50.sb", "98.35.Eg", "Cosmic rays", "Composition energy spectra and interactions", "Electric and magnetic fields", "Astrophysics - High Energy Astrophysical Phenomena" ]
[ "1993A&A...268..726P", "1993A&A...274..902S", "1997JPhG...23..979E", "2001ApJ...547..264J", "2002JHEP...12..033C", "2002PhRvD..66h3004K", "2003A&A...409..799S", "2003JCAP...05..003C", "2003PASJ...55..191N", "2004APh....21..617B", "2004APh....22...47E", "2004NewAR..48..979C", "2005A&A...429..755P", "2005A&A...441..407A", "2005APh....24....1A", "2005JPhG...31R..95H", "2005PhLB..612..147B", "2006PASJ...58..847N", "2006PhRvD..74d3005B", "2007APh....27...76A", "2007JCAP...06..003E", "2008ApJ...678.1165A", "2009ApJ...692L.130A", "2010A&A...522A..67B", "2011APh....35..192G", "2011ApJ...728..122Y", "2011ApJ...738..192P", "2011PhRvL.107q1104A", "2012APh....36..183A", "2012ApJ...746...33A", "2012ApJ...757...14J", "2012ApJ...761L..11J", "2012JCAP...07..031G", "2012MNRAS.426.3295G", "2012PhRvL.108z1101G", "2013A&A...558A..72I", "2013APh....42...15I", "2013APh....47...54A", "2013ApJ...762L..13P", "2013MNRAS.431..415B", "2013MNRAS.436.2326P", "2013PhRvD..87h1101A", "2013PhRvD..88d2004A", "2014BrJPh..44..560L", "2014IAUS..296..188G" ]
[ "10.1103/PhysRevD.90.041302", "10.48550/arXiv.1403.3380" ]
1403
1403.3380_arXiv.txt
14
3
1403.3380
We investigate the possibility that the cosmic ray (CR) knee is entirely explained by the energy-dependent CR leakage from the Milky Way. We test this hypothesis calculating the trajectories of individual CRs with energies between E<SUB>/Z =<SUP>1014</SUP> eV and 10<SUP>17</SUP> eV propagating them in the regular and turbulent Galactic magnetic field. We find a knee-like structure of the CR escape time τ<SUB>esc</SUB>(E) around E/Z =few×<SUP>1015</SUP> eV for a coherence length lc</SUB>≃2 pc of the turbulent field, while the decrease of τ<SUB>esc</SUB>(E) slows down around E/Z ≃<SUP>1016</SUP> eV in models with a weak turbulent magnetic field. Assuming that the injection spectra of CR nuclei are power laws, the resulting CR intensities in such a turbulence are consistent with the energy spectra of CR nuclei determined by KASCADE and KASCADE-Grande. We calculate the resulting CR dipole anisotropy as well as the source rate in this model.
false
[ "E/Z =few×<SUP>1015</SUP", "a weak turbulent magnetic field", "CR nuclei", "models", "SUB>/Z =", "CR", "the turbulent field", "the regular and turbulent Galactic magnetic field", "energies", "τ<SUB>esc</SUB>(E", "KASCADE", "Galactic", "power laws", "the resulting CR intensities", "the resulting CR dipole", "SUP>1014</SUP", "individual CRs", "CR) knee", "the energy-dependent CR leakage", "<" ]
6.240532
0.490095
13
438202
[ "Ragan, S. E.", "Henning, Th.", "Tackenberg, J.", "Beuther, H.", "Johnston, K. G.", "Kainulainen, J.", "Linz, H." ]
2014A&A...568A..73R
[ "Giant molecular filaments in the Milky Way" ]
131
[ "Max Planck Institute for Astronomy, Königstuhl 17, 69117, Heidelberg, Germany", "Max Planck Institute for Astronomy, Königstuhl 17, 69117, Heidelberg, Germany", "Max Planck Institute for Astronomy, Königstuhl 17, 69117, Heidelberg, Germany", "Max Planck Institute for Astronomy, Königstuhl 17, 69117, Heidelberg, Germany", "Max Planck Institute for Astronomy, Königstuhl 17, 69117, Heidelberg, Germany", "Max Planck Institute for Astronomy, Königstuhl 17, 69117, Heidelberg, Germany", "Max Planck Institute for Astronomy, Königstuhl 17, 69117, Heidelberg, Germany" ]
[ "2014A&A...571A..53B", "2014ApJ...797...53G", "2014MNRAS.441.1628S", "2015A&A...584A..67B", "2015AJ....150..159D", "2015ApJ...802...59K", "2015ApJ...805....1B", "2015ApJ...809..154H", "2015ApJ...811..134S", "2015ApJ...815...23Z", "2015EAS....75..105T", "2015MNRAS.447.3390D", "2015MNRAS.450.4043W", "2015PASA...32....7A", "2015esmc.book.....W", "2016A&A...585A.104C", "2016A&A...585A.117Z", "2016A&A...586A..27K", "2016A&A...588A.104G", "2016A&A...590A.131A", "2016A&A...591A...5L", "2016A&A...592A..90C", "2016A&A...595A..32B", "2016ApJ...821...53V", "2016ApJ...832...43O", "2016ApJ...833..246D", "2016ApJS..226....9W", "2016MNRAS.456.2885R", "2016MNRAS.458.3667D", "2016MNRAS.462.3123R", "2016PhDT.......239R", "2016RAA....16...56Z", "2016arXiv160305720L", "2017A&A...598A..76Z", "2017A&A...598A..96L", "2017A&A...600A.141K", "2017A&A...601A.124S", "2017ApJ...835..137W", "2017ApJ...844...15D", "2017ApJ...845...34D", "2017ApJ...848...51D", "2017ApJS..229...24D", "2017ApJS..230....5W", "2017AstRv..13..113V", "2017MNRAS.470.4261D", "2017MNRAS.471.2002L", "2017MNRAS.472.4797S", "2017PASJ...69...78U", "2018A&A...610A..62C", "2018A&A...616A..78M", "2018A&A...618A.159R", "2018A&A...619A.166M", "2018ApJ...852...93V", "2018ApJ...863...59E", "2018ApJ...863...74K", "2018ApJ...864..153Z", "2018MNRAS.475.5268B", "2018MNRAS.479L..40K", "2018arXiv180602286C", "2019A&A...621A.122B", "2019A&A...621A.130L", "2019A&A...622A..52Z", "2019A&A...622A.155Y", "2019A&A...622A.166S", "2019AJ....158..122K", "2019Ap&SS.364..150V", "2019ApJ...880...88X", "2019ApJ...887..186Z", "2019ApJS..245...14E", "2020A&A...634A.139W", "2020A&A...641A..53W", "2020A&A...642A..68S", "2020A&A...642A.163S", "2020ApJ...893...91S", "2020MNRAS.492.1594S", "2020MNRAS.492.5420S", "2020MNRAS.498.5936E", "2020MNRAS.499.3728P", "2020NatAs...4.1064H", "2020RAA....20..143L", "2020RAA....20..159S", "2021A&A...646A.137L", "2021A&A...651L...4S", "2021A&A...655L...2C", "2021ApJ...918L...2L", "2021ApJ...921L..42V", "2021MNRAS.500..191E", "2021MNRAS.500.3027D", "2021MNRAS.500.5268I", "2021RAA....21..188Z", "2021ascl.soft02002W", "2022A&A...657A...1S", "2022A&A...658A..54C", "2022A&A...658A.114K", "2022A&A...658A.128L", "2022A&A...660A..56A", "2022A&A...663A..56N", "2022ApJ...930..169B", "2022ApJ...936..160Z", "2022ApJ...938...44G", "2022ApJS..259...36G", "2022MNRAS.511.2829M", "2022MNRAS.513..638Z", "2022MNRAS.516L..35L", "2022MNRAS.517.5180Z", "2022MNRAS.517L.102L", "2022PASJ...74...24K", "2022RAA....22a5003L", "2022arXiv220807506G", "2023A&A...675A.119G", "2023A&A...676A..27B", "2023A&A...679A.130S", "2023AJ....165...16C", "2023ASPC..534..153H", "2023ASPC..534..233P", "2023ApJ...944..228M", "2023ApJ...944L...8S", "2023ApJ...944L..13T", "2023ApJ...944L..18M", "2023ApJ...949..108K", "2023MNRAS.519.3098C", "2023pcsf.conf..160S", "2024MNRAS.527.3639F", "2024MNRAS.527.4244S", "2024MNRAS.528.1555C", "2024MNRAS.528.2199R", "2024MNRAS.529.1091Z", "2024MNRAS.529.3060G", "2024MNRAS.530.1311F", "2024arXiv240319798E", "2024arXiv240518474Z" ]
[ "astronomy" ]
11
[ "catalogs", "ISM: structure", "Galaxy: structure", "Astrophysics - Astrophysics of Galaxies" ]
[ "1958BAN....14..215W", "1960MNRAS.121..123B", "1979ApJS...41...87S", "1980ApJ...242..528E", "1987ApJ...319..730S", "1987ApJS...63..645U", "1988ApJ...324...60B", "1993ApJ...419L..29E", "1994A&A...291..943O", "1998ApJ...506..306N", "1999ApJ...522..165A", "2001A&A...377.1023L", "2001AJ....122.3017S", "2001ApJ...547..792D", "2001ApJ...553..227P", "2002ApJ...566..945B", "2002ApJ...570..132K", "2003ApJ...596..220C", "2003PASP..115..953B", "2004tra..book.....R", "2005A&A...434..987A", "2005ApJ...634.1126M", "2006ApJ...639..227S", "2006ApJ...643..245V", "2006ApJ...647..997S", "2006ApJS..163..145J", "2006MNRAS.371.1663D", "2006MNRAS.372.1227D", "2007ApJ...665.1138S", "2007MNRAS.375..213W", "2008AJ....135.1301V", "2008ApJ...680..428G", "2008ApJ...689..290H", "2008MNRAS.389.1097D", "2009A&A...493..735L", "2009A&A...497..399K", "2009A&A...504..415S", "2009A&A...505..405P", "2009ApJ...696..484B", "2009ApJ...698..324R", "2009ApJ...699.1153R", "2009ApJ...700..137R", "2009ApJ...700.1609M", "2009ApJ...707..283J", "2010A&A...518L..95H", "2010A&A...518L.102A", "2010A&A...518L.103M", "2010A&A...520A..49S", "2010ApJ...712.1147J", "2010ApJ...719L.185J", "2010ApJ...721..686P", "2010ApJ...723..492R", "2010ApJ...723.1019H", "2010ApJ...724..687L", "2010ApJ...725..534F", "2010ApJ...725.1327S", "2010ApJS..188..123R", "2010PASP..122..314M", "2011A&A...533A..17B", "2011A&A...533A..94H", "2011A&A...536A..48K", "2011ApJS..192....4A", "2011MNRAS.418.1689U", "2012A&A...541A..63P", "2012A&A...542A..77F", "2012A&A...542A.114H", "2012A&A...543L...3H", "2012A&A...544A.146W", "2012ApJ...745..190L", "2012MNRAS.424.2599C", "2012MNRAS.426..701M", "2013A&A...550A.116T", "2013A&A...554A..55H", "2013A&A...556A.153H", "2013A&A...557A.120K", "2013A&A...559A..34L", "2013ApJ...769..115H", "2013ApJ...776...62H", "2013ApJ...779...42S", "2013ApJ...779...44H", "2013ApJS..205....1P", "2013ApJS..208...11L", "2013ApJS..208...14G", "2013ApJS..209....2S", "2013MNRAS.431.1587E", "2014ASSP...36..417B", "2014ApJ...780..173B", "2014ApJ...784....3C", "2014MNRAS.441.1628S" ]
[ "10.1051/0004-6361/201423401", "10.48550/arXiv.1403.1450" ]
1403
1403.1450_arXiv.txt
Star-forming clouds in the Milky Way -- both nearby and distant -- exhibit elongated structures \citep[see][and references therein]{Myers2009}. The morphology appears most enhanced when viewing the densest part of the cloud, which is also the regime most intimately connected to star formation. From the early quiescent phase, filamentary morphology seems to be imprinted on all subsequent stages of star formation. However, the physical origin of filaments is still debated. A variety of models can produce dense, filamentary structures, though unambiguous observational diagnostics are still lacking to determine the dominant mechanism(s) leading to filament formation. Filamentary structures have been observed in a variety of tracers, ranging from extinction maps at optical and infrared wavelengths \citep[e.g.][]{SchneiderElmegreen1979, Apai2005, Jackson2010, Schmalzl2010, Beuther2011, Kainulainen2013} to CO maps \citep{Ungerechts1987, Goldsmith2008, Hacar2013} to far-infrared/sub-millimeter dust emission maps \citep[e.g.][]{Henning2010, Andre2010, Menshchikov2010, Molinari2010, Schneider2010, Hill2011, Hennemann2012, Peretto2012}. The recent results of {\em Herschel} have again highlighted the ubiquity of filaments in the interstellar medium (ISM) and thus rejuvenated interest in the role filamentary morphology plays in star formation. In numerical models, filamentary structure is a natural consequence of a number of dynamic processes in the ISM such as converging flows \citep[e.g.][]{Elmegreen1993, Vazquez-Semadeni2006, Heitsch2008, Clark2012}, the collision of shocked sheets \citep{Padoan2001}, instabilities in self-gravitating sheets \citep[e.g.][]{Nagai1998}, or other analogous processes that compresses gas to an over-dense interface. Such processes commonly occur in a global spiral potential \citep{Dobbs2008}. In non-self-gravitating cases, additional ingredients, such as magnetic fields and/or turbulence dissipation \citep{Hennebelle2013a}, are needed to preserve filaments with the properties observed in the ISM and in turbulent simulations. More massive filaments, such as those modeled by \citet{Fischera2012} and \citet{Heitsch2013a, Heitsch2013b}, are self-gravitating, thus the effects of continuing accretion from large scales and external pressurisation, play an important role in the observed properties. A fundamental quantity in the understanding of the origin of Galactic filaments is the maximum length over which they can occur. This is challenging to observe for several reasons. First, until recently, few unbiased surveys of the Galactic plane that could potentially identify such structures existed. Second, the star formation occurring within filaments, especially massive ones, is very disruptive and impacts the clouds and clear signatures of the parent structure. As such, the quiescent stage of filaments, appearing as infrared-dark clouds (IRDCs), should better preserve the initial formation signatures compared to active clouds. \citet{Jackson2010} identified a long infrared-dark 80\,pc long ``Nessie'' filament. \citet{Goodman2013}\footnote{\url{http://authorea.com/249}} revisited ``Nessie'' finding that it coincides with the Scutum-Centaurus arm, and it may even be at least twice as long. Further searches have found similarly enticing individual structures in the Galactic plane \citep[e.g.][]{Beuther2011,Battersby2012,Tackenberg2013,Li2013}, but to date, no comprehensive compilation of long, coherent structures in the Galaxy exists. In this paper, we present a new sample of Giant Molecular Filaments (GMFs) in the first quadrant of the Milky Way. We describe our methods for identifying filaments in projection using unbiased Galactic plane surveys and our follow-up method for confirming a filament's coherence in velocity space. Our search has produced seven new filaments with lengths on the order of 100\,pc, more than doubling the number of similar structures known from the literature. This catalog will aid in studying the connection between large scale filamentary structure and star formation in the Galaxy. \begin{figure} \includegraphics[width=0.5\textwidth]{267_254_higal_compare.pdf} \caption{\label{fig:higal} Zoom in to the eastern end of the F26.7-25.4 filament. A Grayscale GLIMPSE 8\micron image (top) is plotted with contours at -7.5, -10, -12.5... MJy sr$^{-1}$ (negative to highlight absorption feature), and the HIGAL 250\micron image (bottom) is plotted with contours drawn at 5, 6, 7... Jy beam$^{-1}$.} \end{figure} \begin{figure} \includegraphics[width=0.53\textwidth]{pvdiagram.pdf} \caption{\label{fig:pv} Position-velocity (PV) diagram (integrated over all latitudes) of the region shown in Figure~\ref{fig:higal} based on the Galactic Ring Survey $^{13}$CO(1-0) data. The approximate positions of the spiral arm features in this longitude range from the \cite{Vallee2008} model are labeled in white, and the filament that we identify, GMF26.7-25.4, is labeled in red.} \end{figure}
14
3
1403.1450
Throughout the Milky Way, molecular clouds typically appear filamentary, and mounting evidence indicates that this morphology plays an important role in star formation. What is not known is to what extent the dense filaments most closely associated with star formation are connected to the surrounding diffuse clouds up to arbitrarily large scales. How are these cradles of star formation linked to the Milky Way's spiral structure? Using archival Galactic plane survey data, we have used multiple datasets in search of large-scale, velocity-coherent filaments in the Galactic plane. In this paper, we present our methods employed to identify coherent filamentary structures first in extinction and confirmed using Galactic Ring Survey data. We present a sample of seven giant molecular filaments (GMFs) that have lengths on the order of ~100 pc, total masses of 10<SUP>4</SUP>-10<SUP>5</SUP> M<SUB>⊙</SUB>, and exhibit velocity coherence over their full length. The GMFs we study appear to be inter-arm clouds and may be the Milky Way analogs to spurs observed in nearby spiral galaxies. We find that between 2 and 12% of the total mass (above ~10<SUP>20</SUP> cm<SUP>-2</SUP>) is "dense" (above 10<SUP>22</SUP> cm<SUP>-2</SUP>), where filaments near spiral arms in the Galactic midplane tend to have higher dense gas mass fractions than those further from the arms. <P />Appendix A is available in electronic form at <A href="http://www.aanda.org/10.1051/0004-6361/201423401/olm">http://www.aanda.org</A>
false
[ "Galactic Ring Survey data", "archival Galactic plane survey data", "Galactic Ring Survey", "spiral arms", "velocity coherence", "higher dense gas mass fractions", "Galactic", "coherent filamentary structures", "~100 pc", "lengths", "star formation", "multiple datasets", "nearby spiral galaxies", "total masses", "filaments", "molecular clouds", "inter-arm clouds", "<", "search", "the Galactic plane" ]
11.566105
9.905035
183
437550
[ "Lazorenko, P. F.", "Sahlmann, J.", "Ségransan, D.", "Martín, E. L.", "Mayor, M.", "Queloz, D.", "Udry, S." ]
2014A&A...565A..21L
[ "Astrometric planet search around southern ultracool dwarfs. II. Astrometric reduction methods and a deep astrometric catalogue" ]
15
[ "Main Astronomical Observatory, National Academy of Sciences of the Ukraine, Zabolotnogo 27, 03680, Kyiv, Ukraine", "European Space Agency, European Space Astronomy Centre, P.O. Box 78, Villanueva de la Cañada, 28691, Madrid, Spain; Observatoire de Genève, Université de Genève, 51 Chemin Des Maillettes, 1290, Versoix, Switzerland", "Observatoire de Genève, Université de Genève, 51 Chemin Des Maillettes, 1290, Versoix, Switzerland", "INTA-CSIC Centro de Astrobiología, 28850 Torrejón de Ardoz, Madrid, Spain", "Observatoire de Genève, Université de Genève, 51 Chemin Des Maillettes, 1290, Versoix, Switzerland", "Observatoire de Genève, Université de Genève, 51 Chemin Des Maillettes, 1290, Versoix, Switzerland; University of Cambridge, Cavendish Laboratory, J J Thomson Avenue, Cambridge, CB3 0HE, UK", "Observatoire de Genève, Université de Genève, 51 Chemin Des Maillettes, 1290, Versoix, Switzerland" ]
[ "2014A&A...565A..20S", "2014MmSAI..85..674S", "2015A&A...577A..15S", "2015A&A...579A..61S", "2015MNRAS.453L.103S", "2016A&A...595A..77S", "2016MNRAS.455..357S", "2016SPIE.9836E..02R", "2017A&A...606A..90L", "2018A&A...618A.111L", "2018exha.book.....P", "2019A&A...623A..10T", "2019A&A...629A.113L", "2020MNRAS.495.1136S", "2021MNRAS.500.5453S" ]
[ "astronomy" ]
6
[ "astrometry", "techniques: high angular resolution", "atmospheric effects", "parallaxes", "brown dwarfs", "Astrophysics - Solar and Stellar Astrophysics", "Astrophysics - Earth and Planetary Astrophysics", "Astrophysics - Instrumentation and Methods for Astrophysics" ]
[ "1992AJ....103..638M", "1994AN....315..353A", "1996ApJ...465..264P", "1997KPCB...13b..63L", "1997SPIE.2871.1135A", "1997yCat.1239....0E", "1998Msngr..94....1A", "2001A&A...369..339P", "2003A&A...409..523R", "2003AJ....125..984M", "2004A&A...427.1127L", "2004ApJ...617.1323P", "2005PASP..117.1021S", "2005yCat.2263....0D", "2006A&A...449.1271L", "2009A&A...505..903L", "2010A&A...523A..48J", "2010AJ....139.1844B", "2010Msngr.141....5M", "2011A&A...527A..25L", "2011AdSpR..47..356M", "2012Ap&SS.341...31D", "2013A&A...556A.133S", "2013SPIE.8864E..1BS", "2014A&A...565A..20S", "2014ARA&A..52..339R", "2014ApJ...785..161R" ]
[ "10.1051/0004-6361/201323271", "10.48550/arXiv.1403.4619" ]
1403
1403.4619_arXiv.txt
{\label{notes}} Extrasolar planets around stars can be discovered and characterised by a variety of observation techniques \citep{Seager:2011ve}. Different methods give access to different observables of a given exoplanetary system and are subject to different practical limitations. To obtain a complete picture of extrasolar planets, we therefore depend on having access to the widest possible range of observing techniques. Astrometry consists in measuring the photocentre positions of stellar objects and can be used for indirect exoplanet detection by revealing a star's orbital reflex motion \citep{Sozzetti:2005qy}. However, this relies on a measurement accuracy better than one milli-arcsecond (mas) over time-scales of several years, which requires specialised instruments and methods. Dedicated space missions such as \emph{Hipparcos} \citep{ESA:1997vn} and, in particular, its successor \emph{Gaia} \citep{Perryman:2001vn, de-Bruijne:2012kx} can meet these requirements, and instruments on-board \emph{Hubble} have been used for this purpose, too \citep{Benedict:2010ph}. On the ground, the main practical limitation is turbulence in the Earth's atmosphere \citep{Sahlmann_spie}, which can be overcome with the use of large-aperture telescopes, interferometry, or adaptive optics. For observations of dense stellar fields with the optical camera {\small FORS2} of ESO's Very Large Telescope, we have achieved astrometric precisions of 0.05~mas for well-exposed star images at the field centre when the seeing is restricted to the optimal range { (\citealt{Lazorenko2009}, hereafter \citetalias{Lazorenko2009})}. Because this performance is sufficient for exoplanet detection, we initiated an astrometric search targeting ultracool dwarfs { (UCD)} at the M/L transition in 2010, a survey that is described in detail in the first paper in our series { (\citealt{Sahlmann:2013prep}, hereafter \citetalias{JS2013}).} The detection of planetary signals at or close to the noise level requires high-quality astrometric data, that is, precise position measurements ideally free of systematic errors, and adequate precision estimates that correspond to the measurement accuracy. In this paper, we present the latest optimisations of the astrometric methods for the reduction of {\small FORS2} observations described in { \citetalias{Lazorenko2009}. These methods were already successfully used for the detection of a close 28 Jupiter-mass ($M_J$) companion of an L1.5 dwarf \citep{JS2013} and are implemented for the planet search survey \citepalias{JS2013}. } The paper is structured as follows: the observations are described in Sect. \ref{obs}, the astrometric model is outlined in Sect. \ref{mod}, and the single-frame precision of the measurements is characterised in Sect. \ref{fr_pr}. The epoch residuals and precision of these residuals are introduced in Sect. \ref{ep_res}, and in Sect. \ref{se} we demonstrate the method of the detection and reduction of systematic errors. The observed $\chi^2$ statistic for the epoch residuals and its comparison with the theoretical $\chi^2$ distribution law is analysed in Sect. \ref{conc}. In Sect. \ref{cat}, we describe the catalogue of astrometric and photometric data obtained for field stars in the observed fields. We conclude in Sect. \ref{c_d}.
{\label{c_d}} The {\small FORS2/VLT} observations made for our planet search programme {necessitated} a detailed analysis of the systematic errors in the epoch residuals. We analysed three types of errors: the instability in the relative position of the two CCD chips, systematic errors correlated in space, and errors not correlated in space. The chip motion instability and space-dependent errors are relatively well modelled and suppressed to $\sim$0.05~mas. The residual error $\varphi$ of uncertain origin, however, could not be removed. We only found its r.m.s. value and added it quadratically to the model precision. This component largely determines the single-epoch precision for bright field and dwarf stars, which in the field centre is 0.12--0.17~mas with a median of 0.15~mas. The median r.m.s. of the epoch residuals is a factor $\theta$ smaller because of the noise absorption by the least squares fitting and is 0.126~mas (0.10--0.15~mas for different fields). These estimates refer to the median seeing of 0.6\arcsec and 32 exposures in one epoch. The single-epoch precision for bright stars depends on the field and varies by a factor of two from the best (densely populated but not crowded field Nr.\,20) to typical (Nr.\,5, which is \dwfive), and the worst-case (relatively poor populated star field Nr.\,2) precision. This is shown in Fig. \ref{gaia}, which presents the nominal precision $\sigma_{e}/{\theta_{e}}$ corresponding to an infinite number of epochs. The {vertical}-strip structure of the data points in the plot (each strip corresponds to a separate epoch) is due to the difference in epoch seeing, which varies from 0.4 to 0.9\arcsec. The epochs with best seeing of about 0.40--0.45\arcsec are shown separately by open circles for each field where they form the lower bound with the best precision. In these cases, the nominal {\small FORS2} precision is 0.07--0.15~mas at $I=$16--17~mag. This demonstrates that the nominal epoch precision improves almost proportionally with seeing, and good seeing is favourable for astrometry. In these epochs, however, bright reference star images were often saturated and, even with a single saturated pixel, rejected (Sect. \ref{obs}). The number of suitable images within these epochs decreases, and the epoch precision degrades. Moreover, the rejection of equally bright nearby reference stars increases the reference star noise, which deteriorates the precision even more. The result is an upward trend in precision, most distinct for stars brighter than $I=$17~mag in field Nr.\,2. For field Nr.\,20, the images of selected bright stars never saturate because of stable seeing conditions and the precision improves with brightness. The precision floor due to systematic errors is set by $\phi^2_e$, which is the squared sum of $\varphi$ and the space-correlated error $\phi_{\rm space}(e)$ (Eq. (\ref{eq:upd})). After averaging over all epochs and stars, we obtained the limiting precision $\langle \phi_e \rangle$ marked by the horizontal dashed line in Fig. \ref{gaia}. However, for stars with small space-correlated errors, the precision is better, and the same is true when the small-scale corrections are not measurable due to insufficient number of nearby reference stars and therefore were set to zero. \begin{figure}[!] \centering \resizebox{\hsize}{!}{\includegraphics[bb = 55 49 265 171,clip]{5_pi.eps}} \caption {Distribution of distances, with 1-$\sigma$ uncertainty ranges, and $I$-band magnitudes of catalogue stars with well-measured parallaxes. Theoretical curves correspond to M5, M0, and K0 main-sequence stars. } \label{5pi} \end{figure} The current 0.15~mas estimate of the nominal epoch precision for bright field stars and UCDs specifies the astrometric precision of {\small VLT/FORS2} routine observations under variable observing conditions, with a restriction of 0.9\arcsec\ on seeing alone. Not surprisingly, this precision exceeds the 0.05~mas value obtained in \citetalias{Lazorenko2009} for observations in optimal conditions. Here, because of problems with saturation (Sect. \ref{satur}), we restricted the exposure duration, and therefore the light flux for a single image of a bright star is 0.5--0.7$\times 10^6$ photoelectrons, while the peak flux in the pilot study in \citetalias{Lazorenko2009} was 1.5--2.0$\times 10^6$ photoelectrons. Thus, the label 'bright star' really refers to objects with threefold difference in brightness between both studies. For this reason, the median uncertainty of the photocentre determination increases from 0.23~mas in \citetalias{Lazorenko2009} to 0.49~mas in the current programme, and the reference frame noise increases respectively from 0.15~mas to 0.30~mas. With this two-fold difference, the pilot study estimates scaled to the current observing conditions leads to 0.10~mas epoch precision, in agreement with Fig.\,16 in \citetalias{Lazorenko2009}, and which is similar to the 0.13~mas precision for observations restricted to the optimal seeing (Sect.\ref{av_conc}). We compared the astrometric performance of the VLT and the \emph{Gaia} predictions, which illuminates effects related to differences in the entrance apertures and consequently in the light collecting power. Astrometry of bright objects is surely much better from dedicated space satellites, but the precision for very faint stars favours the use of large ground-based facilities. We illustrated this by comparing the \emph{Gaia} single-epoch precision \citep{GAIA,de-Bruijne:2012kx} with the precision for VLT/{\small FORS2} catalogue stars shown in Fig. \ref{gaia}, given by {the estimator $\sigma_{e}/{\theta_{e}}$ that does not depend on the number of epochs}. For bright stars, \emph{Gaia} can achieve $\sim$0.01~mas single-epoch precision, while at the faint end of $G=20$, this value is expected to be $\sim$0.7--1.1~mas. With the relation $G$$-$$I$ $\sim$ 0.8$-$1.8, valid for spectral classes F8\ldots L2 \citep{Jordi2010, Jordi2010yCat}, this limit corresponds to the range $I=18.2-19.2$ where {\small FORS2} provides us with an epoch precision of 0.2--0.3~mas. For UCDs with $I$-band magnitudes of 16--18, the {\small FORS2} precision is thus up to five times better than the expected \emph{Gaia} precision. This can be reformulated differently: for faint stars, the expected \emph{Gaia} single-measurement precision can be reached with {\small FORS2} for stars that are approximately 4 magnitudes fainter. \begin{figure}[htb] \centering \resizebox{\hsize}{!}{\includegraphics*[bb = 56 61 261 160]{gaiadw2.eps}} \\ \resizebox{\hsize}{!}{\includegraphics*[bb = 56 61 261 160]{gaiadw5.eps}} \\ \resizebox{\hsize}{!}{\includegraphics*[bb = 56 51 261 160]{gaiadw20.eps}} \\ \caption {Nominal single-epoch precision $\sigma_{e}/{\theta_{e}}$ as a function of $I$-band magnitude for catalogue stars in fields Nr. 2, 5, and 20 at any seeing (dots) and for best-seeing epoch (open circles). The target magnitude is marked by an arrow and the precision floor set by systematic errors $\langle \phi_e \rangle$ is shown with a dashed line. } \label{gaia} \end{figure} The main motivations of this study was to provide data with well-predicted statistics, which is critically needed for the detection of low-mass companions to UCDs. We {achieved this goal}, because the obtained epoch residuals are distributed according to a normal law with a scatter parameter equal to the model value $\sigma_{e}$. In addition, the $\chi^2$-values for the epoch residuals corrected with the factor $ c_{\chi^2}^2(I) $ follow the theoretical $\chi^2$-statistic. { The immediate scientific result of {\small FORS2} high precision astrometry is the discovery of two tight binaries, which is discussed in \citetalias{JS2013}, and the measurements of trigonometric parallaxes of 20 UCDs and hundreds of $I =$16-17.5 stars with precision of $\sim 0.1$~mas. This is an excellent performance for ground-based optical astrometry and comes close to the best precisions obtained in the optical with {HST/WFC3}-scanning \citep{HST} and for radio sources with VLBI \citep{VLBI}.} \begin{appendix}{\label{A}}
14
3
1403.4619
<BR /> Aims: We describe the astrometric reduction of images obtained with the FORS2/VLT camera in the framework of an astrometric planet search around 20 M/L-transition dwarfs. We present the correction of systematic errors, the achieved astrometric performance, and a new astrometric catalogue containing the faint reference stars in 20 fields located close to the Galactic plane. <BR /> Methods: Remote reference stars were used both to determine the astrometric trajectories of the nearby planet search targets and to identify and correct systematic errors. <BR /> Results: We detected three types of systematic errors in the FORS2 astrometry: the relative motion of the camera's two CCD chips, errors that are correlated in space, and an error contribution of as yet unexplained origin. The relative CCD motion probably has a thermal origin and typically is 0.001-0.010 px (~0.1-1 mas), but sometimes amounts to 0.02-0.05 px (3-6 mas). This instability and space-correlated errors are detected and mitigated using reference stars. The third component of unknown origin has an amplitude of 0.03-0.14 mas and is independent of the observing conditions. We find that a consecutive sequence of 32 images of a well-exposed star over 40 min at 0.6'' seeing results in a median rms of the epoch residuals of 0.126 mas. Overall, the epoch residuals are distributed according to a normal law with a χ<SUP>2</SUP> value near unity. We compiled a catalogue of 12 000 stars with I-band magnitudes of 16-22 located in 20 fields, each covering ~ 2' × 2'. It contains I-band magnitudes, ICRF positions with 40-70 mas precision, and relative proper motions and absolute trigonometric parallaxes with a precision of 0.1 mas/yr and 0.1 mas at the bright end, respectively. <BR /> Conclusions: This work shows that an astrometric accuracy of ~100 micro-arcseconds over two years can be achieved with a large optical telescope in a survey covering several targets and varying observing conditions. <P />Based on observations made with ESO telescopes at the La Silla Paranal Observatory under programme IDs 086.C-0680, 087.C-0567, 088.C-0679, 089.C-0397, and 090.C-0786.Appendices are available in electronic form at <A href="http://www.aanda.org/10.1051/0004-6361/201323271/olm">http://www.aanda.org</A>The catalogue is available at the CDS via anonymous ftp to <A href="http://cdsarc.u-strasbg.fr">http://cdsarc.u-strasbg.fr</A> (ftp://130.79.128.5) or via <A href="http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/565/A21">http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/565/A21</A>
false
[ "varying observing conditions", "systematic errors", "relative proper motions", "reference stars", "errors", "Remote reference stars", "anonymous ftp", "several targets", "unknown origin", "electronic form", "unity", "~100 micro-arcseconds", "programme IDs", "ESO telescopes", "trigonometric parallaxes", "the observing conditions", "an astrometric planet search", "a new astrometric catalogue", "the nearby planet search targets", "0.1 mas" ]
8.042122
10.307703
-1
437273
[ "Filippova, E.", "Bozzo, E.", "Ferrigno, C." ]
2014A&A...563A.124F
[ "INTEGRAL and Swift observations of the hard X-ray transient MAXI J1828-249" ]
5
[ "ISDC, University of Geneva, Ch. d'Ecogia 16, 1290, Versoix, Switzerland ; Space Research Institute, 117997, Moscow, Russia", "ISDC, University of Geneva, Ch. d'Ecogia 16, 1290, Versoix, Switzerland", "ISDC, University of Geneva, Ch. d'Ecogia 16, 1290, Versoix, Switzerland" ]
[ "2016A&A...587A..61C", "2017MNRAS.471.3494Z", "2019AstL...45..127B", "2019PASJ...71..108O", "2021NewAR..9301618M" ]
[ "astronomy" ]
7
[ "X-rays: binaries", "X-rays: bursts", "accretion", "accretion disks", "black hole physics", "Astrophysics - High Energy Astrophysical Phenomena" ]
[ "1990ARA&A..28..215D", "2003A&A...411L..53C", "2003A&A...411L.131U", "2003A&A...411L.141L", "2003A&A...411L.231L", "2004MNRAS.355.1105F", "2005Ap&SS.300..107H", "2005SSRv..120..165B", "2006A&A...456..917R", "2009MNRAS.398.1194C", "2010LNP...794...53B", "2012A&A...537L...7F", "2013ATel.5474....1N", "2013ATel.5476....1F", "2013ATel.5478....1K", "2013ATel.5479....1K", "2013ATel.5482....1R", "2013ATel.5483....1N", "2013ATel.5484....1M", "2014ATel.5886....1T", "2014ATel.5911....1C" ]
[ "10.1051/0004-6361/201323337", "10.48550/arXiv.1403.1819" ]
1403
1403.1819_arXiv.txt
\label{sec:intro} \maxi\ was discovered in outburst on 15 October 2013 (56580.91 MJD) by MAXI/GSC \citep{nakahira2013}. Shortly after the discovery, the source was detected at higher energies by the hard X-ray imager ISGRI on-board INTEGRAL \citep{filippova2013}. At discovery the flux recorded from the source was 93$\pm$9~mCrab in the 4-10~keV energy band and 45$\pm$2~mCrab (48$\pm$2~mCrab) in the 20-40~keV (40-80~keV) energy band. Follow-up observations carried out on 56581.6 MJD with the narrow field instrument on-board \swift/XRT provided the best measured position of the source at RA=277.2427, Dec=-25.0304 (J2000) with an estimated uncertainty of 3.6~arcsec at 90 c.l. \citep{kennea2013}. The improved position permitted us to identify the UV and IR counterparts of the source. The estimated magnitude of the UV source was 18.64$\pm$0.04 (stat)$\pm$0.03 (sys) \citep[in AB system, without correction for interstellar reddening;][]{kennea2013a}. The AB magnitudes of the optical/IR source were $g' = 17.2\pm0.1, r' = 16.9\pm0.1, i' = 16.9\pm0.1, z' = 16.8\pm0.1, J = 16.8\pm0.1, H = 16.9\pm 0.1, K = 17.2\pm0.2$ (not corrected for the interstellar reddening). The derived spectrum of optical/IR source might be interpreted as arising because of an accretion disk \citep{rau2013}. Radio observations carried out about two days after the onset of the outburst did not reveal significant emission from the source \citep[the 3~$\sigma$ upper limit on its flux was 75~$\mu$Jy at 5.5~GHz and 57~$\mu$Jy at 9.0~GHz;][]{millerjones2013}. In this paper we report on all \inte\ and Swift data available to our group collected during the outburst of \maxi\ from 56580 to 56607~MJD. \begin{figure} \centerline{\includegraphics[scale=0.3]{figures/isgri_mosaics20_100keV.pdf}} \caption{ISGRI mosaic of the FOV around \maxi\ (20-100 keV). The mosaic has been obtained from data collected during revolution 1344.} \label{fig:ima} \end{figure}
\label{sec:discussion} \maxi\ was discovered in outburst on 2013 October 15 by the MAXI/GSC instrument and apparently underwent a spectral transition from a hard to a softer state about one day later \citep{negoro2013}. Such transitions are typically observed in the so-called black hole candidates (BHC), which suggests that \maxi\ is associated with this class of objects. According to the standard scenario \citep{fender2004,homan2005,belloni2010}, BHC sources are known to evolve during their outburst along a q-shaped track in the hardness-intensity diagram. The outburst starts in the low-hard spectral state (LHS), which is characterized by a power-law shaped X-ray spectrum with $\Gamma\sim1.6-1.7$ and a cut-off at the higher energies $E_{cut} \sim 100$ keV. Radio emission arises in this state due to the synchrotron radiation of a steady jet. In the following phases of the outburst, the X-ray and radio luminosities both increase until the source reaches the high-soft state (HSS), characterized by a prominent thermal emission from the accretion disk and a marginal power-law tail. Radio emission at this state is no longer observed, most likely because of the suppression of the jet. The time variability of the source is also significantly different in the two states, being generally more pronounced in the LHS with rms values of up to $\sim$30\% and characterized by the quasi-periodical oscillations (QPOs) in the HSS in combination with a strongly suppressed rms \citep[see, e.g.,][for a recent review]{belloni2010}. In addition to these two main states, at least two intermediate hard and soft states were identified. In particular, BHCs in the so-called intermediate soft state still show a significant hard component (extending to hundreds of keV) though they are no longer significantly detected in the radio domain and present both a prominent soft spectral component and a limited rms. Not all transient BHCs in outburst go through a complete q-track. So far, a limited number of objects were observed that reach the hard intermediate state during outburst, but then return to the LHS instead of moving to HSS \citep[see, e.g.,][]{capitanio2009,ferrigno2012}. The observational campaign we presented in this paper started about 1.8~days after the discovery, that is, immediately after the possible spectral transition reported by \cite{negoro2013}. The broad-band energy spectrum of the source as measured by \swift\ and \inte\ comprised a thermal component that we associated with the emission from a multi-temperature black-body disk, and a hard power-law extending with no measurable break up to 200~keV. Assuming a distance to the source of 8~kpc, its peak luminosity is $\sim10^{38}$ erg/s. During the entire observational campaign \maxi\ did not show significant spectral evolution. Our observational results suggest that \maxi\, is a transient BHC that evolved rapidly in the first two days of the outburst from the LHS to the soft intermediate state and remained in this state during the entire observational campaign presented in this work. This conclusion is supported by both the spectral and timing analysis\footnote{Note that the range of frequencies we could investigate in the case of \maxi\ based on the available X-ray data is slightly lower (0.001-0.5 Hz) compared to that in which typical QPOs for this state are observed (0.01 - 50 Hz). However, the measured rms level in the XRT data is fully compatible with that commonly found in the soft intermediate state.} reported in Sect.~\ref{sec:results} and the lack of significant radio emission already two days after the onset of the outburst, as mentioned in Sect.~\ref{sec:intro}. The light curve of the source based on MAXI publicly available data also supports this scenario and suggests that the source is now fading into quiescence. This means that \maxi\, might represent another example of transient BHC ''failed'' outburst, as mentioned above. Future observations in mid-2014 (e.g. with {\it Swift}/XRT) will be able to confirm/reject this conclusion. {\bf Note added to proofs.} After this paper was accepted for publication, new \swift\, data collected on 2014 February 14 found that the source is back in a faint hard state (power-law photon index $\sim$1.7 and 0.5-10 keV flux of $4\times 10^{-11}$ erg/cm$^2$/s; \cite{tomsick2014}). Radio observations performed on 2014 February 16 also detected significant radio emission from the source (preliminary flux densities 1.38$\pm$0.05 mJy at 5.5 GHz and 1.28$\pm$0.06 mJy at 9 GHz; \cite{corbel2014}), which supports our conclusions in Sect. 5.
14
3
1403.1819
In this paper we report on the observations performed with INTEGRAL and Swift of the first outburst detected from the hard X-ray transient MAXI J1828-249. During the first about two days of the outburst, the source was observed by MAXI to undergo a very rapid transition from a hard to a softer spectral state. While the hard state was not efficiently monitored because the transition occurred so rapidly, the evolution of the source outburst in the softer state was covered quasi-simultaneously in a broad energy range (0.6-150 keV) by the instruments on board INTEGRAL and Swift. During these observations, the spectra measured from the source displayed both a prominent thermal emission with temperature kT ~ 0.7 keV and a power-law hard component with a photon index Γ ~ 2.2 extending to 200 keV. The properties of the source in the X-ray domain are reminiscent of those displayed by black hole transients during the soft intermediate state, which supports the association of MAXI J1828-249 with this class of objects.
false
[ "MAXI J1828", "board INTEGRAL", "objects", "the hard X-ray transient MAXI J1828", "black hole transients", "MAXI", "Swift", "INTEGRAL", "a softer spectral state", "the soft intermediate state", "the softer state", "first", "the hard state", "the X-ray domain", "a broad energy range", "the first outburst", "kT", "temperature", "a photon index Γ", "a power-law hard component" ]
6.194149
6.604844
48
746697
[ "Smith, R. J." ]
2014MNRAS.443L..69S
[ "Variations in the initial mass function in early-type galaxies: a critical comparison between dynamical and spectroscopic results." ]
96
[ "Department of Physics, University of Durham, South Road, Durham DH1 3LE, UK" ]
[ "2014ApJ...792L..37M", "2014ApJ...793...96S", "2014ApJ...795..158M", "2014MNRAS.442.1003P", "2014MNRAS.442.3138F", "2015A&A...573A.110G", "2015ApJ...806L..31M", "2015ApJ...814...26N", "2015IAUS..311...49M", "2015IAUS..311...69Y", "2015MNRAS.446..493P", "2015MNRAS.447.1033M", "2015MNRAS.447.1638M", "2015MNRAS.448L..82F", "2015MNRAS.449.3441S", "2015MNRAS.449.4091C", "2015MNRAS.449L.137L", "2015MNRAS.451.1081M", "2015MNRAS.452..597Z", "2015MNRAS.454.1811B", "2015MNRAS.454L..71S", "2015Msngr.160...18S", "2016ARA&A..54..597C", "2016ASSL..435..585D", "2016ApJ...824...17K", "2016ApJ...831..134V", "2016ApJ...832..203Z", "2016IAUFM..29B.193F", "2016IAUS..315E..19F", "2016IAUS..319...35N", "2016MNRAS.455..308T", "2016MNRAS.456L.104M", "2016MNRAS.457.1468L", "2016MNRAS.459.3677L", "2016MNRAS.462.1415C", "2016MNRAS.462.2832C", "2016MNRAS.462.3854L", "2016MNRAS.463.3220L", "2017ApJ...835..183C", "2017ApJ...837..166C", "2017ApJ...841...32Z", "2017ApJ...841...68V", "2017ApJ...843...74O", "2017ApJ...845..157N", "2017ApJ...846..166L", "2017ApJ...850L..14V", "2017MNRAS.464..453D", "2017MNRAS.464.3812F", "2017MNRAS.467..836S", "2017MNRAS.468..319E", "2017MNRAS.468.1594A", "2017MNRAS.468.3071N", "2017MNRAS.469.2464A", "2017MNRAS.470..815S", "2017MNRAS.471..383S", "2017PhDT........11N", "2018A&A...616A.121S", "2018ApJ...855...20G", "2018ApJ...863...38G", "2018MNRAS.475.1073V", "2018MNRAS.475.2467F", "2018MNRAS.477.3954P", "2018MNRAS.479.2443V", "2018MNRAS.479.5448B", "2018PASA...35...39H", "2019A&A...626A.124M", "2019ApJ...875..151M", "2019ApJ...885...75J", "2019MNRAS.482.2515B", "2019MNRAS.485.4852G", "2019MNRAS.488.1679P", "2019MNRAS.488.5300C", "2019MNRAS.489.5612D", "2020ARA&A..58..577S", "2020FrASS...7...21M", "2020MNRAS.494.5619S", "2020MNRAS.498.4051D", "2021ApJ...923...43V", "2021FrASS...8..157D", "2021MNRAS.500.4153L", "2022A&A...662A...8M", "2022A&A...664A..26P", "2022ApJ...932..103G", "2022MNRAS.509.1959S", "2022MNRAS.514.3410H", "2022MNRAS.516.5712T", "2022PASA...39....2N", "2023A&A...674A..76P", "2023MNRAS.522.5320D", "2023MNRAS.526.4004C", "2024A&A...684A.110M", "2024A&A...686A.122A", "2024ApJ...961..127M", "2024MNRAS.528.7338P", "2024MNRAS.530.3278D", "2024MNRAS.530.4474L" ]
[ "astronomy" ]
5
[ "galaxies: elliptical and lenticular", "cD", "galaxies: stellar content", "Astrophysics - Astrophysics of Galaxies" ]
[ "1955ApJ...121..161S", "2000MNRAS.315..679C", "2001MNRAS.322..231K", "2004ApJ...612..168P", "2009MNRAS.398.1835S", "2010ApJ...709.1195T", "2010ApJ...724..511A", "2010Natur.468..940V", "2012ApJ...747...69C", "2012ApJ...753L..32S", "2012ApJ...760...71C", "2012MNRAS.426.2994S", "2012Natur.484..485C", "2013MNRAS.432.1862C", "2013MNRAS.433.3017L", "2013MNRAS.434.1964S", "2013MNRAS.436..253B", "2014MNRAS.442.1003P", "2015MNRAS.447.1033M" ]
[ "10.1093/mnrasl/slu082", "10.48550/arXiv.1403.6114" ]
1403
1403.6114_arXiv.txt
\begin{figure*} \includegraphics[angle=270,width=175mm]{mlcomp.eps} \vskip -1mm \caption{Comparison of the mass-to-light ratios, $\Upsilon$, used by CvD and A3D, and the derived mismatch parameters, $\alpha$. Panel ``a'' compares the reference mass-to-light ratios, i.e. those derived from stellar population fitting assuming a standard (Kroupa 2001) IMF. Panel ``b'' compares the derived spectroscopic and dynamical mass-to-light ratios. Panel ``c" compares the ratios $\alpha$\,=\,$\Upsilon/\Upsilon_{\rm ref}$, which indicate the inferred deviations from the standard IMF. In panel ``d'', the A3D value of $\alpha$ is modified by using the same $\Upsilon_{\rm ref}$ as used by CvD, for greater consistency (cyan lines show the effect of this change). The red cross in panels ``c'' and ``d'' shows the median value for the two datasets. The quoted slopes and correlation coefficients $r$ are for comparisons in $\log\Upsilon$ or $\log\alpha$; $p$ is the probability of a larger $r$ under the null hypothesis of no correlation. In panels ``a'' and ``b'' the solid line and yellow shading show a linear fit with errors. While the mass-to-light ratios are well correlated between the studies, the ratio $\alpha$ shows essentially no correlation on a galaxy by galaxy basis.} \label{fig:mlcomp} \end{figure*} The stellar initial mass function (IMF) is a key quantity in astrophysics, both intrinsically, as a constraint on the physics of star formation, and also for its importance in converting observed galaxy luminosities into physically-meaningful stellar masses, star-formation rates, etc. In recent years several largely-independent methods have found evidence for a different IMF in early-type galaxies, compared to the Milky Way (MW), and for systematic variation among elliptical galaxies as a function of their mass (Treu et al. 2010; Auger et al. 2010; van Dokkum \& Conroy 2010; Conroy \& van Dokkum 2012b; Smith, Lucey \& Carter 2012; Spiniello et al. 2012; La Barbera et al. 2013; Cappellari et al. 2013). One method to constrain the IMF uses measurements of spectral features that are sensitive to surface gravity at fixed stellar temperature, and hence reveal the presence of low-mass stars in integrated-light spectra (Conroy \& van Dokkum 2012a). Dwarf-star-sensitive features are found to increase in strength in higher mass galaxies, beyond what is expected from element abundance trends, according to spectral synthesis models. This behaviour is interpreted as due to an increasingly bottom-heavy IMF in higher mass galaxies. As noted by Conroy \& van Dokkum (2012b), the signature of bottom-heavy IMFs appears more strongly correlated with Mg/Fe abundance ratios than with velocity dispersion, suggesting that IMF could be linked to star-formation intensity, with more low-mass stars being produced in rapid bursts. Another technique is to infer the total mass in galaxies from gravitational tracers, such as stellar dynamics (e.g. Cappellari et al. 2013) or strong lensing (e.g. Treu et al. 2010). After accounting for the dark matter halo contribution, this leads to an estimate of stellar mass-to-light ratio $\Upsilon$. Comparing this to the ``reference'' mass-to-light ($\Upsilon_{\rm ref}$) expected from the spectrum of the galaxy assuming a MW-like IMF, this yields a mismatch factor $\alpha$\,=\,$\Upsilon/\Upsilon_{\rm ref}$. Lensing and dynamical studies both find a trend of increasing $\alpha$ with galaxy mass, which can be attributed to an increasing contribution of low-mass stars, i.e. a more bottom-heavy IMF. It should be stressed that these two approaches to constraining the IMF measure fundamentally different quantities: gravitational tracers strictly measure mass (which could be dominated by very low-mass dwarfs, or by remnants from massive stars), while the spectroscopic method is sensitive only to low-mass stars. When mass-to-light ratios are quoted from spectroscopy, as by CvD, the values depend on a model assumed for the shape of the IMF. Since spectroscopy and dynamics/lensing each measures a differently-weighted integral over the IMF, comparing the two methods yields a test for the correctness and universality of the assumed IMF model, as well as a test for the systematic errors inherent to each method. For example La Barbera et al. (2013) emphasise that although single- and broken-power-law IMFs can fit their spectroscopic data equally well, the best fitting single-power-law model can be excluded, since it would imply an excessively high $\Upsilon$ for the most massive galaxies. At face value, the recent spectroscopic and dynamical/lensing results do appear to agree, at least at a qualitative level: massive early-type galaxies have IMFs which, on average, are more bottom-heavy than that of the MW, and there is a trend of increasing deviation from the MW IMF at larger mass. This apparent consensus between largely-independent methods has understandably led to increased confidence in these results\footnote{But note that some other works have favoured MW-like IMFs even in very massive ellipticals, using lensing (Smith \& Lucey 2013) and dynamics (J. Thomas et al. in preparation, see http://www.mpa-garching.mpg.de/halo2013/pdfs/day5/11\_thomas.pdf).}. In this Letter, I present a critical evaluation of results obtained by Conroy \& van Dokkum (2012b) and by Cappellari et al. (2013) (hereafter CvD and A3D), where comparisons can be made for exactly the same set of galaxies. I start by directly comparing the spectroscopic and dynamical mass-to-light ratios, the reference mass-to-light ratios, and the IMF mismatch factors, on a galaxy-by-galaxy basis (Section~\ref{sec:mlcomp}). Section~\ref{sec:imftrends} investigates the systematic correlations of mismatch factor with velocity dispersion and Mg/Fe ratios. In Section~\ref{sec:disc}, I highlight the very different systematic trends obtained from CvD and A3D for this common sample of galaxies and discuss possible resolutions. Brief conclusions are drawn in Section~\ref{sec:concs}. \begin{figure*} \includegraphics[angle=0,width=175mm]{alphasig.eps} \vskip -3mm \caption{The relationship between IMF mismatch factor $\alpha$ and velocity dispersion $\sigma$, as derived from spectroscopy by CvD and from dynamics by A3D, for the same sample of galaxies. Panel ``c'' shows A3D adjusted to use the CvD-derived reference mass-to-light ratio, for consistency. Slopes and correlation coefficients are for $\log\alpha$ versus $\log\sigma$. The solid line and yellow shading show the linear fit with errors to the data in each panel.} \label{fig:alphasig} \end{figure*} \begin{figure*} \includegraphics[angle=0,width=175mm]{alphamgfe.eps} \vskip -3mm \caption{Equivalent to Figure~\ref{fig:alphasig} but now for correlations of IMF mismatch factor $\alpha$ with abundance ratio [Mg/Fe]. Slopes and correlation coefficients are for $\log\alpha$ versus [Mg/Fe]. Note how, compared to the previous figure, the correlation is strengthened for CvD but weakened for A3D. } \label{fig:alphamgfe} \end{figure*}
\label{sec:concs} I have presented a comparison between spectroscopic (CvD) and dynamical (A3D) results on the IMF in elliptical galaxies, using a common sample of 34 galaxies with measurements in both studies. The analysis shows that ``consensus'' between dynamical and spectroscopic measurements is present only at the most rudimentary level: both approaches find that a heavier-than-MW IMW is required, on average, for the common sample, but there is no correlation between the mass-excess factors derived from the two methods, on a galaxy-by-galaxy level. The two studies apparently find a correlation of $\alpha$ with some quantity related to galaxy mass. When plotted only against $\sigma$, there is reasonable agreement in slope, but this treatment obscures a clear discrepancy between the results: the correlation found by CvD is not a trend with $\sigma$, but entirely with the Mg/Fe abundance ratio. By contrast, A3D finds no correlation of $\alpha$ with Mg/Fe and is hence in significant conflict with the spectroscopic method. The sense of this disagreement could indicate that confounding factors such as dark matter contributions (A3D) or unusual abundance patterns (CvD) have not been correctly separated from the IMF effects in one or other of the methods. Alternatively, since the two methods are sensitive to different aspects of the IMF, further comparison between dynamical and spectroscopic estimates of $\alpha$ might lead to a more detailed understanding of the {\it shape} of the IMF and its possible variation in elliptical galaxies. Work is ongoing to derive spectroscopic IMF constraints for more galaxies in the A3D sample, and to improve the treatment of element abundance treatment in the CvD models (C. Conroy, private communication). Hence, an enlarged and updated comparison between the two methods should be possible in the near future.
14
3
1403.6114
I present a comparison between published dynamical (ATLAS3D) and spectroscopic (Conroy &amp; van Dokkum) constraints on the stellar initial mass function (IMF) in early-type galaxies, using the 34 galaxies in common between the two works. Both studies infer an average IMF mass factor α (the stellar mass relative to a Kroupa-IMF population of similar age and metallicity) greater than unity, i.e. both methods favour an IMF which is heavier than that of the Milky Way, on average over the sample. However, on a galaxy-by-galaxy basis, there is no correlation between α inferred from the two approaches. I investigate how the two estimates of α are correlated systematically with the galaxy velocity dispersion, σ, and with the Mg/Fe abundance ratio. The spectroscopic method, based on the strengths of metal absorption lines, yields a correlation only with metal abundance ratios: at fixed Mg/Fe, there is no residual correlation with σ. The dynamical method, applied to exactly the same galaxy sample, yields the opposite result: the IMF variation correlates only with dynamics, with no residual correlation with Mg/Fe after controlling for σ. Hence, although both methods indicate a heavy IMF on average in ellipticals, they lead to incompatible results for the systematic trends, when applied to the same set of galaxies. The sense of the disagreement could suggest that one (or both) of the methods has not accounted fully for the main confounding factors, i.e. element abundance ratios or dark matter contributions. Alternatively, the poor agreement might indicate additional variation in the detailed shape of the IMF, beyond what can currently be inferred from the spectroscopic features.
false
[ "metal abundance ratios", "IMF", "galaxies", "galaxy", "dark matter contributions", "an average IMF mass factor", "incompatible results", "metal absorption lines", "fixed Mg/Fe", "similar age", "the Mg/Fe abundance ratio", "α", "i.e. element abundance ratios", "additional variation", "early-type galaxies", "the stellar initial mass function", "the IMF variation", "Mg/Fe", "the galaxy velocity dispersion", "Kroupa-IMF" ]
10.59402
8.453099
126
745667
[ "Henshaw, J. D.", "Caselli, P.", "Fontani, F.", "Jiménez-Serra, I.", "Tan, J. C." ]
2014MNRAS.440.2860H
[ "The dynamical properties of dense filaments in the infrared dark cloud G035.39-00.33" ]
109
[ "School of Physics and Astronomy, University of Leeds, Leeds LS2 9JT, UK", "School of Physics and Astronomy, University of Leeds, Leeds LS2 9JT, UK", "INAF-Osservatorio Astrofisico di Arcetri, L.go E. Fermi 5, I-50125 Firenze, Italy", "European Southern Observatory, Karl-Schwarzschild-Str. 2, D-85748 Garching, Germany", "Department of Astronomy, University of Florida, Gainesville, FL 32611, USA" ]
[ "2015A&A...573A.119R", "2015A&A...584A..67B", "2015AJ....150..159D", "2015ApJ...798...75B", "2015ApJ...802...75K", "2015ApJ...805....1B", "2015ApJ...809..154H", "2015ApJ...811...56W", "2015EAS....75..237M", "2015EAS....75..245P", "2015MNRAS.452.2410S", "2015esmc.book.....W", "2016A&A...587A..69W", "2016A&A...587A..96P", "2016A&A...587A.118P", "2016A&A...595A.122L", "2016ApJ...818...59D", "2016ApJ...822...59S", "2016ApJ...827..107P", "2016ApJ...832..186S", "2016IAUS..315...95T", "2016IAUS..315..146M", "2016MNRAS.456.2586H", "2016MNRAS.457.2675H", "2016MNRAS.458.1139P", "2016MNRAS.458.1990B", "2016MNRAS.459.1803W", "2016MNRAS.463..146H", "2016MNRAS.463L.122H", "2017A&A...598A..76Z", "2017A&A...600A.141K", "2017A&A...601A..18B", "2017A&A...601A.124S", "2017A&A...602L...2H", "2017A&A...603A..89K", "2017A&A...606A.133S", "2017ApJ...841...88W", "2017ApJ...841...97S", "2017IAUS..316....9M", "2017MNRAS.464L..31H", "2017RAA....17...35L", "2018A&A...610A..77H", "2018A&A...611L...3S", "2018A&A...618L...5N", "2018A&A...619A.166M", "2018ARA&A..56...41M", "2018ApJ...855....9L", "2018ApJ...859..151L", "2018MNRAS.474.3760C", "2018MNRAS.475.5268B", "2018MNRAS.478L..54B", "2019A&A...621A.130L", "2019A&A...628A..78R", "2019A&A...629A..81T", "2019A&A...632A..83S", "2019ApJ...872...30S", "2019ApJ...875...24C", "2019ApJ...886..102S", "2019MNRAS.487.1259L", "2020ApJ...894...87C", "2020ApJ...897...53R", "2020ApJ...905...78Y", "2020IAUS..345...51H", "2020MNRAS.491.4310T", "2020MNRAS.493.3412F", "2020MNRAS.497..793Z", "2020MNRAS.499.1666C", "2020NatAs...4.1064H", "2020SSRv..216...76B", "2021A&A...649A.113B", "2021A&A...650A.202R", "2021ApJ...908...86A", "2021MNRAS.503.4320F", "2021MNRAS.503.4601B", "2021MNRAS.505.5183R", "2022A&A...662A...8M", "2022A&A...665A..22B", "2022A&A...667A...6C", "2022ApJ...926..165L", "2022ApJ...927..106C", "2022ApJ...936..169R", "2022ApJ...938...44G", "2022MNRAS.513.1244G", "2022MNRAS.514.6038Z", "2022MNRAS.515.2266A", "2022MNRAS.517.5272H", "2022arXiv220807506G", "2023A&A...674A..75N", "2023A&A...675A..53B", "2023A&A...676A..69Z", "2023ASPC..534....1C", "2023ASPC..534..153H", "2023ASPC..534..233P", "2023ApJ...950..148M", "2023ApJ...955..104R", "2023MNRAS.519.2391Z", "2023MNRAS.519.3851P", "2023MNRAS.520.2306T", "2023MNRAS.525.2935P", "2023RAA....23a5019S", "2024A&A...681A..51W", "2024A&A...686A.146Z", "2024ApJ...961..176B", "2024MNRAS.528.1172R", "2024MNRAS.528.6370F", "2024arXiv240300917K", "2024arXiv240308091H", "2024arXiv240407363A", "2024arXiv240415862Z" ]
[ "astronomy" ]
16
[ "stars: formation", "ISM: clouds", "ISM: individual objects: G035.39-00.33", "ISM: kinematics and dynamics", "ISM: molecules", "Astrophysics - Solar and Stellar Astrophysics", "Astrophysics - Galaxy Astrophysics" ]
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[ "10.1093/mnras/stu446", "10.48550/arXiv.1403.1444" ]
1403
1403.1444_arXiv.txt
\label{Section:introduction} Understanding and categorising the initially quiescent phases of massive ($>$\,8\,\sol) star formation is \emph{essential} if we are to develop a complete picture of how massive stars form. Once star formation is under way, the disruptive effect of stellar feedback destroys the primordial information needed to explain their formation. Consequently, the search for massive starless cores (the dense precursors to massive stars), requires the identification of relatively \emph{quiescent} clouds, that have yet to be affected by feedback from massive young stellar objects. Infrared dark clouds (hereafter, IRDCs), were discovered as extended structures, silhouetted against the bright mid-infrared (MIR) emission observed towards the Galactic centre \citep{perault_1996, egan_1998}. IRDCs are categorised as having large-masses ($\sim$\,10$^{2-5}$\,\sol; \citealp{rathborne_2006, jouni_2013}, hereafter KT13), high-column densities ($\rm N_{H_2}$\,$\sim$\,10$^{22-25}$\,cm$^{-2}$; \citealp{egan_1998, carey_1998}), and low temperatures ($\leq$\,20\,K; \citealp{pillai_2006, peretto_2010, ragan_2011, chira_2013}), making them ideal environments to study the \emph{initial conditions} of star formation. In addition, the high mass surface densities ($\sim$\,0.3\,g\,cm$^{-2}$; \citealp{butler_2012}, hereafter BT12) and high volume densities ($\sim$\,10$^{4-6}$\,cm$^{-3}$; \citealp{carey_1998, rathborne_2006}, BT12) of IRDC clumps are most akin to regions of massive star formation \citep{tan_2013}. This study is the sixth instalment of a series of papers whose goal is to provide a detailed case study of the chemistry, dynamics, and physical structure of a single IRDC, \irdc. \citet{butler_2009} (hereafter, BT09) selected 10 IRDCs (due to their high-contrast against the Galactic MIR background) from the sample of 38 IRDCs studied by \citet{rathborne_2006}. From this sample, \irdc, cloud H of BT09, was chosen for further study because: i) it has one of the most extreme filamentary morphologies in the \citet{rathborne_2006} study; ii) it exhibits extended quiescent regions with little or no signatures of star formation activity (4.5\,\micron, 8\,\micron, 24\,\micron \ emission; \citealp{chambers_2009, carey_2009}); iii) it is relatively nearby, with a kinematic distance of $\sim$ 2900\,pc \citep{simon_2006}. Since 2010, \irdc \ has been revealed to be an extremely complex, globally virialised structure (Paper III; \citealp{hernandez_2012a}), consisting of several, morphologically distinct filaments (Paper IV; \citealp{henshaw_2013}) exhibiting common velocity gradients (Paper V; \citealp{izaskun_2014}), that is in an early stage of evolution (Paper II; \citealp{hernandez_2011}). Paper I \citep{izaskun_2010}, discovered the presence of faint and narrow SiO emission ($<$\,1\,\kms), traced over parsec scales in \irdc. Two main suggestions were put forward to explain this emission: i) a population of widespread, and undetected low-mass protostars (the IRAM\,30\,m beam at $\sim$\,87\,GHz is $\sim$\,28\arcsec); ii) the emission may be a large-scale shock product of the cloud formation process. Whilst a population of deeply-embedded protostars cannot be ruled out without higher-angular resolution observations of shocked gas tracers, the
\label{Section:conclusions} We have presented a detailed kinematic study using high-sensitivity and high-spectral resolution PdBI observations of \ntwoh \ (1-0) towards a highly-filamentary IRDC. Our results and analysis lead us to conclude the following: \begin{enumerate} \item Multiple filaments are identified both spectrally and spatially. F2a, F2b, and F3 have mean centroid velocities of (45.34\,$\pm$\,0.04)\,\kms, (46.00\,$\pm$\,0.05)\,\kms, (46.86\,$\pm$\,0.04)\,\kms, respectively. \item The abrupt change in velocity noted at the location of H6 (Papers IV \& V), rather than being indicative of large scale flows towards H6, may be explained by the presence of substructure within filament 2, i.e. F2a and F2b. \item F2a, F2b, and F3 have mean line-widths (FWHM) of (0.83\,$\pm$\,0.04)\,\kms, (0.77\,$\pm$\,0.04)\,\kms, and (0.71\,$\pm$\,0.04)\,\kms, respectively. The ratio of non-thermal to thermal (for \ntwoh) velocity dispersion for each velocity component is 5.4, 5.0, and 4.7, respectively. The ratio of the non-thermal component of the line-width to the isothermal sound speed for an average molecule (mass = 2.33\,a.m.u.) at 15\,K are 1.6, 1.4, and 1.4, respectively. This indicates that the gas motions are \emph{mildly} supersonic. In regions where multiple spectral components are evident, moment analysis can overestimate the non-thermal contribution to the line-width by a factor $\gtrsim$\,2. \item Globally, the kinematics of the gas are relatively quiescent, indicated by the small velocity gradients observed over each filament (of the order $<$\,0.7\vel). Locally, however, the mean velocity gradients can reach $\sim$ 1.5--2.5\,\vel. \item There is some indication that the kinematics of the dense gas may be influenced by the self-gravity of dense cores within filaments, or possibly by outflow feedback from already forming stars. Further molecular line observations are required to discern between these two scenarios. For these two opposing scenarios we calculate: \begin{enumerate} \item \emph{Infall}: The mass accretion rate is estimated to be $\sim$\,(7\,$\pm$4)$\times$10$^{-5}$\,\sol\,yr$^{-1}$. The filaments retain their structure within the vicinity of H6, and individual filaments appear to feed individual cores. The SW continuum core could accrete an additional (36\,$\pm$\,25)\,\sol, in an estimated free-fall time of (5\,$\pm$\,3)\,$\times$10$^{5}$\,yrs. \item \emph{Expanding shell}: The momentum for the expanding shell is estimated to be $\sim$\,(24\,$\pm$\,12)\,\sol\,\kms. The dense filamentary structures may have been separated from the main body of IRDC material due to the dynamic processes of star formation. \end{enumerate} \end{enumerate} \noindent Our analysis highlights the importance of combining high-sensitivity and high-spectral resolution data at high-angular resolution, to put quantitative constraints on the dynamics of high-mass star forming regions.
14
3
1403.1444
Infrared dark clouds (IRDCs) are unique laboratories to study the initial conditions of high-mass star and star cluster formation. We present high-sensitivity and high-angular-resolution Institut de Radioastronomie Millimétrique (IRAM) Plateau de Bure Interferometer observations of N<SUB>2</SUB>H<SUP>+</SUP> (1-0) towards IRDC G035.39-00.33. It is found that G035.39-00.33 is a highly complex environment, consisting of several mildly supersonic filaments (σ _NT/c<SUB>s</SUB> ∼ 1.5), separated in velocity by &lt;1 km s<SUP>-1</SUP>. Where multiple spectral components are evident, moment analysis overestimates the non-thermal contribution to the line-width by a factor of ∼2. Large-scale velocity gradients evident in previous single-dish maps may be explained by the presence of substructure now evident in the interferometric maps. Whilst global velocity gradients are small (&lt;0.7 km s<SUP>-1</SUP> pc<SUP>-1</SUP>), there is evidence for dynamic processes on local scales (∼1.5-2.5 km s<SUP>-1</SUP> pc<SUP>-1</SUP>). Systematic trends in velocity gradient are observed towards several continuum peaks. This suggests that the kinematics are influenced by dense (and in some cases, starless) cores. These trends are interpreted as either infalling material, with accretion rates ∼(7 ± 4) × 10<SUP>-5</SUP> M<SUB>⊙</SUB> yr<SUP>-1</SUP>, or expanding shells with momentum ∼24 ± 12 M<SUB>⊙</SUB> km s<SUP>-1</SUP>. These observations highlight the importance of high-sensitivity and high-spectral-resolution data in disentangling the complex kinematic and physical structure of massive star-forming regions.
false
[ "lt;1 km s", "dynamic processes", "local scales", "pc", "evidence", "velocity gradient", "Plateau de Bure Interferometer observations", "Plateau de Bure Interferometer", "global velocity gradients", "IRDC G035.39", "several continuum peaks", "Institut de Radioastronomie Millimétrique", "substructure", "∼2", "velocity", "Large-scale velocity gradients", "∼(7 ±", "previous single-dish maps", "small (&lt;0.7 km s", "massive star-forming regions" ]
11.371617
10.540079
183
1154458
[ "Amorín, R.", "Pérez-Montero, E.", "Contini, T.", "Vílchez, J. M.", "Bolzonella, M.", "Tasca, L. A. M.", "Lamareille, F.", "Zamorani, G.", "Maier, C.", "Carollo, C. M.", "Kneib, J. -P.", "Le Fèvre, O.", "Lilly, S.", "Mainieri, V.", "Renzini, A.", "Scodeggio, M.", "Bardelli, S.", "Bongiorno, A.", "Caputi, K.", "Cucciati, O.", "de la Torre, S.", "de Ravel, L.", "Franzetti, P.", "Garilli, B.", "Iovino, A.", "Kampczyk, P.", "Knobel, C.", "Kovač, K.", "Le Borgne, J. -F.", "Le Brun, V.", "Mignoli, M.", "Pellò, R.", "Peng, Y.", "Presotto, V.", "Ricciardelli, E.", "Silverman, J. D.", "Tanaka, M.", "Tresse, L.", "Vergani, D.", "Zucca, E." ]
2015A&A...578A.105A
[ "Extreme emission-line galaxies out to z ~ 1 in zCOSMOS. I. Sample and characterization of global properties" ]
96
[ "INAF-Osservatorio Astronomico di Roma, via Frascati 33, 00040 Monteporzio Catone, Roma, Italy ; Instituto de Astrofísica de Andalucía, CSIC, 18008, Granada, Spain", "Instituto de Astrofísica de Andalucía, CSIC, 18008, Granada, Spain; IRAP, Université de Toulouse, UPS-OMP, Toulouse, France; Institut de Recherche en Astrophysique et Planétologie, CNRS, 14 avenue Édouard Belin, 31400, Toulouse, France", "IRAP, Université de Toulouse, UPS-OMP, Toulouse, France; Institut de Recherche en Astrophysique et Planétologie, CNRS, 14 avenue Édouard Belin, 31400, Toulouse, France", "Instituto de Astrofísica de Andalucía, CSIC, 18008, Granada, Spain", "INAF-Osservatorio Astronomico di Bologna, via Ranzani 1, 40127, Bologna, Italy", "Laboratoire d'Astrophysique de Marseille, CNRS-Université d'Aix-Marseille, 38 rue Frederic Joliot Curie, 13388, Marseille, France", "IRAP, Université de Toulouse, UPS-OMP, Toulouse, France; Institut de Recherche en Astrophysique et Planétologie, CNRS, 14 avenue Édouard Belin, 31400, Toulouse, France", "INAF-Osservatorio Astronomico di Bologna, via Ranzani 1, 40127, Bologna, Italy", "Institute of Astronomy, ETH Zurich, 8093, Zürich, Switzerland; University of Vienna, Department of Astronomy, Tuerkenschanzstrasse 17, 1180, Vienna, Austria", "Institute of Astronomy, ETH Zurich, 8093, Zürich, Switzerland", "Laboratoire d'Astrophysique de Marseille, CNRS-Université d'Aix-Marseille, 38 rue Frederic Joliot Curie, 13388, Marseille, France", "Laboratoire d'Astrophysique de Marseille, CNRS-Université d'Aix-Marseille, 38 rue Frederic Joliot Curie, 13388, Marseille, France", "Institute of Astronomy, ETH Zurich, 8093, Zürich, Switzerland", "European Southern Observatory, Karl-Schwarzschild-Strasse 2, 85748, Garching b. Muenchen, Germany", "Dipartimento di Astronomia, Universitá di Padova, vicolo Osservatorio 3, 35122, Padova, Italy", "INAF-IASF Milano, via Bassini 15, 20133, Milano, Italy", "INAF-Osservatorio Astronomico di Bologna, via Ranzani 1, 40127, Bologna, Italy", "Max-Planck-Institut für extraterrestrische Physik, 85748, Garching b. Muenchen, Germany", "Kapteyn Astronomical Institute, University of Groningen, 9700 AV, Groningen, The Netherlands", "INAF-Osservatorio Astronomico di Trieste, via Tiepolo, 11, 34143, Trieste, Italy", "SUPA Institute for Astronomy, The University of Edinburgh, Royal Observatory, Edinburgh, EH9 3HJ, UK", "SUPA Institute for Astronomy, The University of Edinburgh, Royal Observatory, Edinburgh, EH9 3HJ, UK", "INAF-IASF Milano, via Bassini 15, 20133, Milano, Italy", "INAF-IASF Milano, via Bassini 15, 20133, Milano, Italy; Laboratoire d'Astrophysique de Marseille, CNRS-Université d'Aix-Marseille, 38 rue Frederic Joliot Curie, 13388, Marseille, France", "INAF-Osservatorio Astronomico di Brera, via Brera, 28, 20159, Milano, Italy", "Institute of Astronomy, ETH Zurich, 8093, Zürich, Switzerland", "Institute of Astronomy, ETH Zurich, 8093, Zürich, Switzerland", "Institute of Astronomy, ETH Zurich, 8093, Zürich, Switzerland; MPA - Max-Planck Institut für Astrophysik, Karl-Schwarzschild-Str. 1, 85741, Garching, Germany", "IRAP, Université de Toulouse, UPS-OMP, Toulouse, France; Institut de Recherche en Astrophysique et Planétologie, CNRS, 14 avenue Édouard Belin, 31400, Toulouse, France", "Laboratoire d'Astrophysique de Marseille, CNRS-Université d'Aix-Marseille, 38 rue Frederic Joliot Curie, 13388, Marseille, France", "INAF-Osservatorio Astronomico di Bologna, via Ranzani 1, 40127, Bologna, Italy", "IRAP, Université de Toulouse, UPS-OMP, Toulouse, France; Institut de Recherche en Astrophysique et Planétologie, CNRS, 14 avenue Édouard Belin, 31400, Toulouse, France", "Institute of Astronomy, ETH Zurich, 8093, Zürich, Switzerland", "Universitá degli Studi dell'Insubria, via Valleggio 11, 22100, Como, Italy; INAF-Osservatorio Astronomico di Brera, via Brera, 28, 20159, Milano, Italy", "Instituto de Astrofísica de Canarias, vía Lactea s/n, 38200, La Laguna, Tenerife, Spain", "IPMU, Institute for the Physics and Mathematics of the Universe, 5-1-5 Kashiwanoha, 38200, 277-8583x, Kashiwa, Japan", "IPMU, Institute for the Physics and Mathematics of the Universe, 5-1-5 Kashiwanoha, 38200, 277-8583x, Kashiwa, Japan", "Laboratoire d'Astrophysique de Marseille, CNRS-Université d'Aix-Marseille, 38 rue Frederic Joliot Curie, 13388, Marseille, France", "INAF-Osservatorio Astronomico di Bologna, via Ranzani 1, 40127, Bologna, Italy; INAF-IASFBO, via P. Gobetti 101, 40129, Bologna, Italy", "INAF-Osservatorio Astronomico di Bologna, via Ranzani 1, 40127, Bologna, Italy" ]
[ "2014A&A...568L...8A", "2014ApJ...788L...4A", "2014ApJ...791...17M", "2014MNRAS.445.3200S", "2014arXiv1412.3853P", "2015A&A...576A.116V", "2015ApJ...799...36R", "2015ApJ...800...18A", "2015ApJ...805...45L", "2015MNRAS.449..917L", "2015PASJ...67..104T", "2015sf2a.conf...77D", "2016A&A...585A..51D", "2016A&A...588A...5T", "2016A&A...588A..23A", "2016A&A...588A..58B", "2016A&A...594A..49G", "2016A&A...596A..67R", "2016ApJ...828...67L", "2016MNRAS.455.3359P", "2016MNRAS.458..449M", "2016MNRAS.459.2591G", "2016PhyU...59....3V", "2016arXiv160107368T", "2017A&A...599A..28K", "2017A&A...601A..95C", "2017A&A...605A..42C", "2017ApJ...834..181O", "2017ApJ...842...47V", "2017ApJ...843L...8B", "2017MNRAS.467.1287P", "2017MNRAS.469.4705C", "2017MNRAS.472.2608S", "2017NatAs...1E..52A", "2017PASP..129d3001P", "2018ApJ...854...29M", "2018ApJ...869..141C", "2018MNRAS.473.1258S", "2018MNRAS.476.3793K", "2018MNRAS.478..791N", "2018PASJ...70S..35M", "2018cagn.conf...49P", "2019A&A...621A..52L", "2019A&A...629A...9A", "2019A&ARv..27....3M", "2019ApJ...872..145J", "2019MNRAS.489.1787B", "2020ApJ...898...45T", "2020ApJ...904...59A", "2020ApJ...904..180O", "2020IAUS..352...84S", "2020MNRAS.491.2254W", "2020MNRAS.494.3541H", "2021ApJ...909..130R", "2021MNRAS.502.6044E", "2021MNRAS.503.4105T", "2021MNRAS.504.1237P", "2021MNRAS.505..903C", "2022A&A...659A..16L", "2022A&A...663A..29B", "2022A&A...664A.110B", "2022A&A...665A..95I", "2022A&A...668A..60L", "2022ApJ...928...78R", "2022ApJ...935..150L", "2022ApJ...940L..52B", "2022MNRAS.511.2515F", "2022MNRAS.513.4451B", "2022MNRAS.513.5211T", "2022MNRAS.514.2627C", "2022MNRAS.516..975T", "2022PASJ...74....1T", "2023A&A...671A..99E", "2023A&A...676A..53L", "2023ApJ...945...35T", "2023ApJ...952...76M", "2023ApJ...953...10C", "2023ApJ...953L..29L", "2023ApJ...957...39L", "2023MNRAS.518..592K", "2023MNRAS.522.2089D", "2023MNRAS.522.2655G", "2023MNRAS.523..270R", "2023MNRAS.525.2087B", "2023arXiv231207799D", "2024A&A...686A..60C", "2024ApJ...963...54P", "2024ApJ...967...73J", "2024ApJ...967L..17S", "2024MNRAS.528.3340B", "2024arXiv240116934B", "2024arXiv240305362L", "2024arXiv240312683C", "2024arXiv240410811M", "2024arXiv240417415D", "2024arXiv240616725T" ]
[ "astronomy" ]
52
[ "galaxies: evolution", "galaxies: fundamental parameters", "galaxies: abundances", "galaxies: starburst", "galaxies: irregular", "galaxies: star formation", "Astrophysics - Astrophysics of Galaxies", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
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[ "10.1051/0004-6361/201322786", "10.48550/arXiv.1403.3441" ]
1403
1403.3441_arXiv.txt
Low-mass galaxies undergoing vigorous bursts of star formation over galaxy-wide scales provide unique laboratories for understanding {galaxy} mass assembly and chemical evolution over cosmic times. {In the local Universe, these systems are often referred to as H{\sc II} galaxies \citep{Terlevich1991} and Blue Compact Dwarfs \citep[BCDs;][]{Thuan1981}, depending on the observational technique or the selection criteria \citep[see][for a review]{KunthOstlin2000}. In spectroscopic surveys, they are generally recognized by their high-excitation emission lines with unusually large equivalent widths (EW)\footnote{We use the convention of positive equivalent widths for emission lines.}, as a product of the photoionization of gas by hot massive stars in a young burst of star formation \citep{Sargent1970}.} Over the last decade, the advent of all-sky optical and UV surveys such as the {\it Sloan Digital Sky Survey} \citep[SDSS;][]{Abazajian2003} and the {\it Galaxy Evolution Explorer} \citep[GALEX;][]{Martin2005}, along with other smaller surveys, have allowed us to systematically search and characterize relatively large samples of extreme emission-line galaxies (EELGs) out to the frontiers of the local Universe \citep[$z \la 0.3$, e.g.,][]{Kniazev2004,Kakazu2007,Overzier2008,Salzer2009,Cardamone2009, Cowie2010,Izotov2011,Shim2013}. {This has made} it possible to discover an increasing number of extremely compact, low-metallicity galaxies with unusually high {specific} star formation rates (SFR, sSFR$=$SFR/M$_{\star} \sim$1-100\,Gyr$^{-1}$), such as the {\it \emph{green peas}} \citep{Cardamone2009,Amorin2010} and a handful of {extremely metal-poor galaxies \citep[XMPs; $Z$\,$\la$\,0.1\,$Z_{\odot}$][]{KunthOstlin2000}} at $0.1$\,$\la$$z$$\la$\,0.4 \citep[e.g.,][]{Kakazu2007,Hu2009,Cowie2010}. Similarly to nearby H{\sc II} galaxies and some BCDs, EELGs {are probably} the youngest and chemically least evolved population of low$-z$ star-forming galaxies \citep[SFGs, e.g.,][]{Searle1972,Rosa2007,Jaskot2013}. % These properties make them unique probes with which to study the details of chemical enrichment, massive star formation, and feedback processes in galaxies with physical properties (i.e., size, mass, SFR, metallicity, gas, and dust relative content) most closely resembling those prevailing at high redshift, e.g., Lyman-break galaxies and Lyman-$\alpha$ emitters \citep[e.g.,][]{Pettini2001,Finkelstein2011}. {Furthermore, increasing observational evidence point to EELGs as} the likely environments to host the progenitors of long-duration gamma-ray bursts \citep[e.g.,][]{Christensen2004,Kewley2007,Savaglio2009,Guseva2011} and the most luminous supernovae \citep{Chen2013,Lunnan2013,Leloudas2014,Toene2014}. In order to understand comprehensively the properties of EELGs as a class, to select best case studies for detailed analysis, and {to} provide a valuable benchmark for comparative studies at higher redshifts, large and representative samples of EELGs must be assembled. {Although} EELGs are generally rare among local low-mass galaxies \citep[$<$0.5\% of galaxies in SDSS;] []{Kniazev2004}, their frequency and significance in the context of the cosmic star formation rate density is expected to increase out to $z\sim$\,1 \citep{Guzman1997,Kakazu2007}. However, because of their faintness and compactness, studying EELGs at these intermediate redshifts requires a great deal of observational effort. Thus, pioneering studies have been limited to relatively small samples of intrinsically luminous objects \citep[e.g.,][]{Koo1995,Phillips1997}. In this context, recent deep multiwavelength surveys have offered a new avenue for studying chemical enrichment and starburst activity and its associated feedback processes in strongly star-forming EELGs out to $z\sim$\,1 and beyond \citep[see, e.g,][]{Hoyos2005,vanderWel2011,Atek2011,Trump2011,Xia2012, Henry2013,Ly2014,Amorin2014a,Amorin2014b,Masters2014,Maseda2014}. This is the case of the COSMOS survey \citep{Scoville2007} and one of its spectroscopic follow-ups, the zCOSMOS 20k bright survey \citep{Lilly2007}. In particular, the wealth of {high-quality photometric and spectroscopic data} provided by these surveys allow {us to perform} a thorough and systematic characterization of a large probe of faint ($I_{\rm AB}\la$\,22.5 mag) EELGs out to $z\sim$\,1. While the large collection of deep broad- and narrowband photometric measurements provided by COSMOS allows luminosities and reliable stellar masses to be derived, HST-ACS $I-$band imaging provides the spatial resolution {required} to study morphological properties. Moreover, zCOSMOS provides the {high signal-to-noise (S/N) spectroscopy required to properly measure the emission lines used to derive reliable gas-phase metallicities.} Remarkably, and despite {the challenge of measuring temperature sensitive emission line ratios (e.g., [\oiii]\,5007/4363), zCOSMOS spectroscopy allows gas-phase metallicity to be derived using the so-called} {\it direct} ($T_{\rm e}$) an unprecedentedly large number of EELGs at intermediate redshifts. Thus, our survey offers the opportunity of identifying a relatively large number of extremely metal-deficient ($\la$\,0.1 $Z_{\odot}$) galaxy candidates. {This is the first of a series of papers aimed at investigating the formation history and evolution of low-mass star-forming galaxies over cosmological time scales using deep multiwavelength surveys. Here, we present the largest spectroscopic sample of galaxies with extreme nebular emission in the range 0.1\,$\la$$z$$\la$\,1 assembled so far. We characterize more than 150 EELGs selected from the zCOSMOS 20k bright survey, based on different key properties, namely size, stellar mass, metallicity, and SFR, which are discussed as a function of morphology and environment. } The {derived properties will be used in a companion paper (Amor\'in et al., in prep.; Paper\,{\sc II}) to discuss possible evolutionary scenarios based on their position in scaling relations involving mass, size, metallicity, and SFR.} Our paper is organized as follows. In Section\,\ref{sect:sample} we describe the parent sample, our dataset, and the selection criteria adopted to {compile} the sample of EELGs. In Section\,\ref{sect:properties} {we present the main physical properties of the sample. We describe the methodology used to derive stellar masses, star formation rates and UV properties, and gas-phase metallicities. As part of the analysis, we also present an alternative method aimed at obtaining $T_{\rm e}-$based metallicities in those EELGs without available measurements of the [\oii]\,3727,3729 doublet. We finish Section\,\ref{sect:properties} by studying the morphological and environmental properties of EELGs.} Later, in {Sections\,\ref{sect:discussion1}-\ref{sect:discussion3}, we highlight the discovery of a number of extremely metal-poor galaxy candidates, discuss the connection between EELGs and Ly$\alpha$ emitters, and compare the zCOSMOS EELGs with other previous samples. Finally, Section\,\ref{sect:conclusions} summarizes our main results and conclusions. } Throughout this paper we adopt the standard $\Lambda$-CDM cosmology, {\em \emph{i.e.}}, $h$ = 0.7, $\Omega_m$ = 0.3, and $\Omega_\Lambda$ = 0.7 (Spergel et al., 2007) {and a solar metallicity value of 12$+\log$(O/H)$=$8.69 \citep{AllendePrieto2001}. Magnitudes are given in the AB system. }
\label{sect:conclusions} {Using the zCOSMOS 20k bright survey we have selected a large sample of 183 extreme emission-line galaxies (EELGs) at $0.11 \leq z \leq 0.93$ showing unusually high [\oiii]$\lambda$5007 rest-frame equivalent widths (EW([\oiii])$\geq$100\AA). We have used zCOSMOS optical spectroscopy and multiwavelength COSMOS photometry and HST-ACS {\it I}-band imaging to characterize the main properties of EELGs, such as sizes, stellar masses, SFR, and metallicity, as well as morphology and large-scale environment. } We summarize our main findings as follows: \begin{enumerate} \item {The adopted} selection criterion based on EW([\oiii]) {lead to a sample of galaxies with the highest EWs in all the observed strong emission lines, e.g., H$\beta$ ($\ga$\,20\AA) and H$\alpha$ ($\ga$\,100\AA), suggesting galaxies dominated by young ($\la$\,10 Myr) star-forming regions}. {The EELGs constitute} 3.4\% of {SFGs} in our parent zCOSMOS sample. Using emission-line diagnostic diagrams we divided the sample {into} 165 purely star-forming galaxies plus 18 NL-AGN candidates ($\sim$\,10\%). {Only four of them are detected as bright X-ray sources.} \item EELGs form the low-end of stellar mass and {the} high-end of sSFR distributions of SFGs in zCOSMOS up to $z\sim$1. Stellar masses of EELGs, as derived from multiband SED fitting, lie in the range $7 \la \log$\,(M$_*$/M$_{\odot}) \la 10$. {Our sample, however, is not complete in mass below $\sim$\,10$^9$\,M$_{\odot}$ in the considered redshift range.} Star formation rates from both \ha\ and FUV luminosities after corrections for dust attenuation and extinction are consistent with each other and range 0.1\,$\la$\,SFR\,$\la$\,35\,M${_\odot}$\,yr$^{-1}$ (Chabrier IMF). Both quantities increase similarly with redshift, so this results {in} a tight range of {\it specific} SFRs (median sSFR$=$\,10$^{-8.18}$\,yr$^{-1}$) and stellar mass doubling times 0.01\,Gyr$<$M$_{*}/$SFR$<$1\,Gyr. \item EELGs are characterized by their low metallicities, 7.3\,$\la$\,12$+\log$(O/H)\,$\la$\,8.5 (0.05-0.6\,$Z_{\odot}$), as derived using both the direct measurements based on electron temperature ($t_{\rm e}$) and strong-line methods calibrated consistently with galaxies with $t_{\rm e}$ measurements. Therefore, the chemical abundances of EELGs at $0.11 \leq z \leq 0.93$ are very similar to those of nearby H{\sc II} galaxies and BCDs. {We find six ($\sim$4\%) extremely metal-poor ($Z <$\,0.1\,$Z_{\odot}$) galaxies in our sample. } \item EELGs are moderately low-dust, very compact UV-luminous galaxies, as evidenced by their typically blue colors ($\beta \sim$\,$-1.6$), high FUV luminosities ($L_{\rm FUV}$\,$\sim$\,10$^{10.4} L_{\odot}$) and high surface brightnesses $\mu_{\rm FUV}$\,$\ga$\,10$^{9} L_{\odot}$\,kpc$^{-2}$. We find only four EELGs with GALEX-UV spectroscopic observations. All these galaxies are strong Ly$\alpha$ emitters, with large equivalent widths and luminosities in the ranges EW(Ly$\alpha$)$=$22-45\AA\ and $\log(L_{\rm Ly\alpha})$\,$=$\,41.8-42.4 erg s$^{-1}$, respectively. \item Using HST-ACS $I-$band COSMOS images we classify star-forming EELGs {into} four morphological classes according to the distribution and shape of their high- and low-surface-brightness components (i.e., SF knots and {underlying} galaxy, respectively). {We} show that 18\% have {round/nucleated} morphologies, {most of which are barely resolved, {while} the remaining 82\% have irregular morphologies. These irregular morphologies are visually classified as} {\it clumpy/chain} (37\%), {\it cometary/tadpole} (16\%), and {\it merger/interacting} (29\%). Therefore, we conclude that at least $\sim$80\% of the EELG sample shows non-axisymmetric morphologies. {Using quantitative morphological parameters} we find that EELGs show smaller half-light radii ($r_{50} \sim$1.3\,kpc in the median) and larger concentration, asymmetry, and Gini parameters than other SFGs in zCOSMOS, most of them being classified as {\it irregular} galaxies {by} automated {algorithms}. {Among the defined morphological classes we do not find any significant difference in the redshift distribution or physical properties. } \item {As star-forming dwarfs in the Local Universe, EELGs are usually found in relative isolation. While only very few EELGs belong to compact groups, almost one third of them are found in spectroscopically confirmed {loose} pairs or triplets. Comparing isolated and grouped EELGs we do not find any significant differences in the redshift distributions or physical properties. } \end{enumerate} In conclusion, we have shown that galaxies selected by their extreme strength of optical emission lines led us to a homogeneous, representative sample of compact, low-mass, low-metallicity, vigorously star-forming systems identifiable with luminous, higher-$z$ versions of nearby H{\sc II} galaxies and blue compact dwarfs. The extreme properties of some of these rare systems closely resemble those of luminous compact galaxies, such as the {\it green peas} \citep{Cardamone2009,Amorin2010} and other samples of emission line galaxies with very high equivalent widths recently found at similar and higher redshift \citep[e.g.,][]{Hoyos2005, Kakazu2007,Salzer2009,Izotov2011,Atek2011,vanderWel2011,vanderWel2013,Xia2012, Shim2013,Henry2013,Ly2014,Maseda2014,Amorin2014b}. The EELGs are galaxies likely caught in a transient and early period of their evolution, where they are efficiently building up a significant fraction of their present-day stellar mass in a young, galaxy-wide starburst. Therefore, they constitute an ideal benchmark for comparative studies with samples of high redshift Ly$\alpha$ emitters and Lyman-break galaxies of similar mass and high-ionization state.
14
3
1403.3441
Context. The study of large and representative samples of low-metallicity star-forming galaxies at different cosmic epochs is of great interest to the detailed understanding of the assembly history and evolution of low-mass galaxies. <BR /> Aims: We present a thorough characterization of a large sample of 183 extreme emission-line galaxies (EELGs) at redshift 0.11 ≤ z ≤ 0.93 selected from the 20k zCOSMOS bright survey because of their unusually large emission line equivalent widths. <BR /> Methods: We use multiwavelength COSMOS photometry, HST-ACS I-band imaging, and optical zCOSMOS spectroscopy to derive the main global properties of star-forming EELGs, such as sizes, stellar masses, star formation rates (SFR), and reliable oxygen abundances using both "direct" and "strong-line" methods. <BR /> Results: The EELGs are extremely compact (r<SUB>50</SUB> ~ 1.3 kpc), low-mass (M<SUB>∗</SUB> ~ 10<SUP>7</SUP>-10<SUP>10</SUP> M<SUB>⊙</SUB>) galaxies forming stars at unusually high specific star formation rates (sSFR ≡ SFR/M<SUB>⋆</SUB> up to 10<SUP>-7</SUP> yr<SUP>-1</SUP>) compared to main sequence star-forming galaxies of the same stellar mass and redshift. At rest-frame UV wavelengths, the EELGs are luminous and show high surface brightness and include strong Lyα emitters, as revealed by GALEX spectroscopy. We show that zCOSMOS EELGs are high-ionization, low-metallicity systems, with median 12+log (O/H) = 8.16 ± 0.21 (0.2 Z<SUB>⊙</SUB>) including a handful of extremely metal-deficient (&lt;0.1 Z<SUB>⊙</SUB>) EELGs. While ~80% of the EELGs show non-axisymmetric morphologies, including clumpy and cometary or tadpole galaxies, we find that ~29% of them show additional low-surface-brightness features, which strongly suggests recent or ongoing interactions. As star-forming dwarfs in the local Universe, EELGs are most often found in relative isolation. While only very few EELGs belong to compact groups, almost one third of them are found in spectroscopically confirmed loose pairs or triplets. <BR /> Conclusions: The zCOSMOS EELGs are galaxies caught in a transient and probably early period of their evolution, where they are efficiently building up a significant fraction of their present-day stellar mass in an ongoing, galaxy-wide starburst. Therefore, the EELGs constitute an ideal benchmark for comparison studies between low- and high-redshift low-mass star-forming galaxies. <P />Full Tables 1 and 2 are only available at the CDS via anonymous ftp to <A href="http://cdsarc.u-strasbg.fr">http://cdsarc.u-strasbg.fr</A> (ftp://130.79.128.5) or via <A href="http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/578/A105">http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/578/A105</A>
false
[ "galaxies", "star formation rates", "stars", "zCOSMOS EELGs", "main sequence star-forming galaxies", "low-mass galaxies", "high surface brightness", "optical zCOSMOS", "low-metallicity star-forming galaxies", "stellar masses", "star-forming EELGs", "unusually high specific star formation rates", "strong Lyα emitters", "non-axisymmetric morphologies", "EELGs", "star-forming dwarfs", "redshift", "low- and high-redshift low-mass star-forming galaxies", "different cosmic epochs", "reliable oxygen abundances" ]
12.699312
6.903131
165
483369
[ "Park, H.", "Han, C.", "Gould, A.", "Udalski, A.", "Sumi, T.", "Fouqué, P.", "Choi, J. -Y.", "Christie, G.", "Depoy, D. L.", "Dong, Subo", "Gaudi, B. S.", "Hwang, K. -H.", "Jung, Y. K.", "Kavka, A.", "Lee, C. -U.", "Monard, L. A. G.", "Natusch, T.", "Ngan, H.", "Pogge, R. W.", "Shin, I. -G.", "Yee, J. C.", "μFUN Collaboration", "Szymański, M. K.", "Kubiak, M.", "Soszyński, I.", "Pietrzyński, G.", "Poleski, R.", "Ulaczyk, K.", "Pietrukowicz, P.", "Kozłowski, S.", "Skowron, J.", "Wyrzykowski, Ł.", "OGLE Collaboration", "Abe, F.", "Bennett, D. P.", "Bond, I. A.", "Botzler, C. S.", "Chote, P.", "Freeman, M.", "Fukui, A.", "Fukunaga, D.", "Harris, P.", "Itow, Y.", "Koshimoto, N.", "Ling, C. H.", "Masuda, K.", "Matsubara, Y.", "Muraki, Y.", "Namba, S.", "Ohnishi, K.", "Rattenbury, N. J.", "Saito, To.", "Sullivan, D. J.", "Sweatman, W. L.", "Suzuki, D.", "Tristram, P. J.", "Wada, K.", "Yamai, N.", "Yock, P. C. M.", "Yonehara, A.", "MOA Collaboration" ]
2014ApJ...787...71P
[ "OGLE-2012-BLG-0455/MOA-2012-BLG-206: Microlensing Event with Ambiguity in Planetary Interpretations Caused by Incomplete Coverage of Planetary Signal" ]
7
[ "Department of Physics, Institute for Astrophysics, Chungbuk National University, Cheongju 371-763, Korea", "Department of Physics, Institute for Astrophysics, Chungbuk National University, Cheongju 371-763, Korea; Author.; The μFUN Collaboration.", "Department of Astronomy, The Ohio State University, 140 West 18th Avenue, Columbus, OH 43210, USA; The μFUN Collaboration.", "Warsaw University Observatory, Al. Ujazdowskie 4, 00-478 Warszawa, Poland; The OGLE Collaboration.", "Department of Earth and Space Science, Osaka University, Osaka 560-0043, Japan; The MOA Collaboration.", "IRAP, CNRS, Université de Toulouse, F-31400 Toulouse, France", "Department of Physics, Institute for Astrophysics, Chungbuk National University, Cheongju 371-763, Korea", "Auckland Observatory, Auckland, New Zealand", "Department of Physics and Astronomy, Texas A&amp;M University, College Station, TX 77843, USA", "Kavli Institute for Astronomy and Astrophysics, Peking University, Yi He Yuan Road 5, Hai Dian District, Beijing 100871, China", "Department of Astronomy, The Ohio State University, 140 West 18th Avenue, Columbus, OH 43210, USA", "Department of Physics, Institute for Astrophysics, Chungbuk National University, Cheongju 371-763, Korea", "Department of Physics, Institute for Astrophysics, Chungbuk National University, Cheongju 371-763, Korea", "Department of Astronomy, The Ohio State University, 140 West 18th Avenue, Columbus, OH 43210, USA", "Korea Astronomy and Space Science Institute, 776 Daedukdae-ro, Yuseong-gu, Daejeon 305-348, Korea", "Kleinkaroo Observatory, Calitzdorp, and Bronberg Observatory, Pretoria, South Africa", "Auckland Observatory, Auckland, New Zealand; Institute for Radiophysics and Space Research, AUT University, Auckland, New Zealand", "Auckland Observatory, Auckland, New Zealand", "Department of Astronomy, The Ohio State University, 140 West 18th Avenue, Columbus, OH 43210, USA", "Department of Physics, Institute for Astrophysics, Chungbuk National University, Cheongju 371-763, Korea", "Department of Astronomy, The Ohio State University, 140 West 18th Avenue, Columbus, OH 43210, USA; Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138, USA; Sagan Fellow.", "-", "Warsaw University Observatory, Al. Ujazdowskie 4, 00-478 Warszawa, Poland", "Warsaw University Observatory, Al. Ujazdowskie 4, 00-478 Warszawa, Poland", "Warsaw University Observatory, Al. Ujazdowskie 4, 00-478 Warszawa, Poland", "Warsaw University Observatory, Al. Ujazdowskie 4, 00-478 Warszawa, Poland; Departamento de Astronomia, Universidad de Concepción, Casilla 160-C, Concepción, Chile", "Warsaw University Observatory, Al. Ujazdowskie 4, 00-478 Warszawa, Poland; Department of Astronomy, The Ohio State University, 140 West 18th Avenue, Columbus, OH 43210, USA", "Warsaw University Observatory, Al. Ujazdowskie 4, 00-478 Warszawa, Poland", "Warsaw University Observatory, Al. Ujazdowskie 4, 00-478 Warszawa, Poland", "Warsaw University Observatory, Al. Ujazdowskie 4, 00-478 Warszawa, Poland", "Warsaw University Observatory, Al. Ujazdowskie 4, 00-478 Warszawa, Poland", "Warsaw University Observatory, Al. Ujazdowskie 4, 00-478 Warszawa, Poland; Institute of Astronomy, University of Cambridge, Madingley Road, Cambridge CB3 0HA, UK", "-", "Solar-Terrestrial Environment Laboratory, Nagoya University, Nagoya 464-8601, Japan", "Department of Physics, University of Notre Dame, 225 Nieuwland Science Hall, Notre Dame, IN 46556-5670, USA", "Institute of Information and Mathematical Sciences, Massey University, Private Bag 102-904, North Shore Mail Centre, Auckland, New Zealand", "Department of Physics, University of Auckland, Private Bag 92-019, Auckland 1001, New Zealand", "School of Chemical and Physical Sciences, Victoria University, Wellington, New Zealand", "Department of Physics, University of Auckland, Private Bag 92-019, Auckland 1001, New Zealand", "Okayama Astrophysical Observatory, National Astronomical Observatory of Japan, Asakuchi, Okayama 719-0232, Japan", "Solar-Terrestrial Environment Laboratory, Nagoya University, Nagoya 464-8601, Japan", "School of Chemical and Physical Sciences, Victoria University, Wellington, New Zealand", "Solar-Terrestrial Environment Laboratory, Nagoya University, Nagoya 464-8601, Japan", "Department of Earth and Space Science, Osaka University, Osaka 560-0043, Japan", "Institute of Information and Mathematical Sciences, Massey University, Private Bag 102-904, North Shore Mail Centre, Auckland, New Zealand", "Solar-Terrestrial Environment Laboratory, Nagoya University, Nagoya 464-8601, Japan", "Solar-Terrestrial Environment Laboratory, Nagoya University, Nagoya 464-8601, Japan", "Solar-Terrestrial Environment Laboratory, Nagoya University, Nagoya 464-8601, Japan", "Department of Earth and Space Science, Osaka University, Osaka 560-0043, Japan", "Nagano National College of Technology, Nagano 381-8550, Japan", "Department of Physics, University of Auckland, Private Bag 92-019, Auckland 1001, New Zealand", "Tokyo Metropolitan College of Aeronautics, Tokyo 116-8523, Japan", "School of Chemical and Physical Sciences, Victoria University, Wellington, New Zealand", "Institute of Information and Mathematical Sciences, Massey University, Private Bag 102-904, North Shore Mail Centre, Auckland, New Zealand", "Department of Earth and Space Science, Osaka University, Osaka 560-0043, Japan", "Mount John University Observatory, PO Box 56, Lake Tekapo 8770, New Zealand", "Department of Earth and Space Science, Osaka University, Osaka 560-0043, Japan", "Department of Physics, Faculty of Science, Kyoto Sangyo University, 603-8555 Kyoto, Japan", "Department of Physics, University of Auckland, Private Bag 92-019, Auckland 1001, New Zealand", "Department of Physics, Faculty of Science, Kyoto Sangyo University, 603-8555 Kyoto, Japan", "-" ]
[ "2016ApJ...820...79B", "2016ApJ...833..145S", "2017ApJ...841...75J", "2018AJ....156..136M", "2018AJ....156..226H", "2019AJ....158..114H", "2021A&A...649A..90H" ]
[ "astronomy" ]
3
[ "gravitational lensing: micro", "planets and satellites: general", "Astrophysics - Earth and Planetary Astrophysics" ]
[ "1986A&A...164..237S", "1986A&A...166...36K", "1991ApJ...374L..37M", "1992ApJ...392..442G", "1992ApJ...396..104G", "1993PASP..105.1342S", "1997MNRAS.284..172W", "1998ApJ...500...37G", "1998ApJ...503..325A", "1998ApJ...506..533G", "1999A&A...349..108D", "2000A&A...363.1081C", "2000AcA....50..421W", "2000ApJ...534..894A", "2001MNRAS.327..868B", "2002ApJ...572..521A", "2003AcA....53..291U", "2003ApJ...591..204S", "2004ApJ...611..528G", "2005ApJ...630..535C", "2006ApJ...644L..37G", "2009AN....330....4T", "2009MNRAS.397.2099A", "2011ApJ...735...85S", "2011ApJ...738...87S", "2011MNRAS.412..607P", "2011Natur.473..349S", "2012ARA&A..50..411G", "2012ApJ...756...48C", "2013ApJ...778...55H", "2013ApJ...778..134P", "2014ApJ...786...85J", "2014MNRAS.439..604S" ]
[ "10.1088/0004-637X/787/1/71", "10.48550/arXiv.1403.1672" ]
1403
1403.1672_arXiv.txt
Gravitational microlensing is one of important methods to detect and characterize extrasolar planets. Due to its sensitivity to planets that are otherwise difficult to detect, the microlensing method is complementary to other methods. In particular, the method is sensitive to planets of low-mass stars located at or beyond the snow line, low-mass planets including terrestrial planets \citep{jung14}, and even free-floating planets \citep{sumi11}. For general review of planetary microlensing, see \citet{gaudi12}. The microlensing signal of a planet is usually a short-term perturbation to the smooth and symmetric standard light curve of the primary-induced lensing event \citep{mao91,gould92b}. The planetary perturbation occurs when the source approaches planet-induced caustics, that represent the positions on the source plane at which the magnification of a point source would become infinite. For a lens composed of a star and a planet, caustics form a single or multiple sets of closed curves each of which is composed of concave curves that meet at cusps. The number, size, and shape of caustics vary depending on the separation and the mass ratio between the planet and its host star. For a given planetary system, planetary perturbations further vary depending on how the source approaches the lens. As a result, planets exhibit very diverse signals in lensing light curves. Due to the immense diversity of planetary signals, characterizing a microlensing planet is a difficult task. This characterization is done from modeling in which an observed lensing light curve is compared to numerous theoretical curves resulting from various combinations of the parameters describing the lens and the source. In this process, it is often confronted that solutions of different lensing parameters result in similar light curves and can explain the observed light curve. This degeneracy problem causes difficulty in the unique interpretation of the lens system. Therefore, understanding the causes of various types of degeneracy is very important. Up to now, it is known that there exist three broad categories of degeneracy in the interpretation of planetary microlensing signals. The first category corresponds to the case for which the degeneracy occurs when different planetary systems induce similar caustics. Good examples are the ``close/wide'' degeneracy for binary-lens events \citep{griest98,dominik99,an05} and the ``ecliptic'' degeneracy for events affected by parallax effects \citep{skowron11}. The second category is that the degeneracy occurs when light curves accidentally appear to be similar despite the fact that the caustics of the degenerate solutions are very different. \citet{choi12} presented two examples of events for which an observed perturbation could be interpreted by either a planetary or a binary companion. The third category is that perturbations can be interpreted by solutions of totally different origins. A good example is the binary-lens/binary-source degeneracy \citep{gaudi98,gaudi04,hwang13}. In this work, we show that incomplete coverage of a perturbation can also result in degenerate solutions even for events where the planetary signal is detected with a high level of statistical significance. We demonstrate the degeneracy for an actually observed event OGLE-2012-BLG-0455/MOA-2012-BLG-206.
We analyzed the high magnification microlensing event OGLE-2012-BLG-0455/MOA-2012-BLG-206 for which the peak of the light curve exhibited anomaly. Despite a large deviation from a standard point-mass model, it was found that there existed 4 very degenerate local solutions. While two of these were due to the well-known $s\leftrightarrow s^{-1}$ ``close/wide'' degeneracy, the other degeneracy, between high and low mass ratios q, was previously unknown. From the fact that the model light curves of the latter degeneracy substantially differed in the parts that were not covered by observation, it was found that the degeneracy was caused by the incomplete coverage of the perturbation. Therefore, the event illustrated the importance of continuous coverage of perturbations for accurate determinations of lens properties. It is expected that the frequency of the degeneracy introduced in this work will be greatly reduced with the improvement of lensing surveys. Recently, there has been such improvement for the existing lensing surveys. For example, The observation cadence of the OGLE lensing survey was substantially increased with the adoption of a new wide field camera. The recent joint of the Wise survey \citep{shvartzvald14} being conducted in Israel enables more continuous event coverage by filling the gap between telescopes in Oceania and Chile. Furthermore, there are plans for future lensing surveys. For example, the MOA group plans to additionally locate a new telescope in Africa for better coverage of lensing events. In addition, the Korea Microlensing Telescope Network (KMTNet) will start operation from the 2014 season by using a network of telescopes at three different locations of the Southern Hemisphere (Chile, South Africa, and Australia). The KMTNet project plans to achieve 10-minute cadence. In addition to survey experiments, there also has been important progress in follow-up experiments. The most important is the completion of Las Cumbres Observatory Global Telescope Network, which is an integrated set of robotic telescopes distributed around the world, including two 2-m telescopes in Hawaii and Australia and nine 1-m telescopes sited in Chile, South Africa, Australia, and Texas \citep{tsapras09}. With the expansion of both survey and follow-up experiments, round-the-clock coverage of lensing events will be possible and the occurrence of the degeneracy will be greatly decreased.
14
3
1403.1672
Characterizing a microlensing planet is done by modeling an observed lensing light curve. In this process, it is often confronted that solutions of different lensing parameters result in similar light curves, causing difficulties in uniquely interpreting the lens system, and thus understanding the causes of different types of degeneracy is important. In this work, we show that incomplete coverage of a planetary perturbation can result in degenerate solutions even for events where the planetary signal is detected with a high level of statistical significance. We demonstrate the degeneracy for an actually observed event OGLE-2012-BLG-0455/MOA-2012-BLG-206. The peak of this high-magnification event (A <SUB>max</SUB> ~ 400) exhibits very strong deviation from a point-lens model with Δχ<SUP>2</SUP> &gt;~ 4000 for data sets with a total of 6963 measurements. From detailed modeling of the light curve, we find that the deviation can be explained by four distinct solutions, i.e., two very different sets of solutions, each with a twofold degeneracy. While the twofold (so-called close/wide) degeneracy is well understood, the degeneracy between the radically different solutions is not previously known. The model light curves of this degeneracy differ substantially in the parts that were not covered by observation, indicating that the degeneracy is caused by the incomplete coverage of the perturbation. It is expected that the frequency of the degeneracy introduced in this work will be greatly reduced with the improvement of the current lensing survey and follow-up experiments and the advent of new surveys.
false
[ "new surveys", "degeneracy", "degenerate solutions", "different lensing parameters", "similar light curves", "solutions", "different types", "statistical significance", "data sets", "incomplete coverage", "the current lensing survey", "events", "an observed lensing light curve", "a twofold degeneracy", "four distinct solutions", "twofold", "the light curve", "the degeneracy", "this degeneracy", "observation" ]
14.310319
12.264888
3
885796
[ "Centeno, R.", "Schou, J.", "Hayashi, K.", "Norton, A.", "Hoeksema, J. T.", "Liu, Y.", "Leka, K. D.", "Barnes, G." ]
2014SoPh..289.3531C
[ "The Helioseismic and Magnetic Imager (HMI) Vector Magnetic Field Pipeline: Optimization of the Spectral Line Inversion Code" ]
93
[ "High Altitude Observatory, NCAR, 3080 Center Green Dr., 80301, Boulder, CO, USA", "HEPL Solar Physics, Stanford University, 94305-4085, Stanford, CA, USA; Max Planck Institute for Solar System Research, Max-Planck-Str. 2, 37191, Katlenburg-Lindau, Germany; Max Planck Institute for Solar System Research, Max-Planck-Str. 2, 37191, Katlenburg-Lindau, Germany", "HEPL Solar Physics, Stanford University, 94305-4085, Stanford, CA, USA", "HEPL Solar Physics, Stanford University, 94305-4085, Stanford, CA, USA", "HEPL Solar Physics, Stanford University, 94305-4085, Stanford, CA, USA", "HEPL Solar Physics, Stanford University, 94305-4085, Stanford, CA, USA", "NwRA, 3380 Mitchell Lane, 80301, Boulder, CO, USA", "NwRA, 3380 Mitchell Lane, 80301, Boulder, CO, USA" ]
[ "2014A&A...572A..54B", "2014ApJ...788L..18S", "2014MNRAS.445...49B", "2014SoPh..289.3231G", "2014SoPh..289.3483H", "2014SoPh..289.3549B", "2014arXiv1405.7353S", "2015A&A...583A..47P", "2015ApJ...798..114S", "2015ApJ...806..245V", "2015ApJ...807..102S", "2015ApJ...808...72S", "2015ApJS..219...17Y", "2015IAUS..305...97P", "2015SoPh..290.1159T", "2015SoPh..290.1507H", "2015SoPh..290.2171H", "2015SoPh..290.2199L", "2016ApJ...823..101S", "2016ApJ...828...12V", "2016ApJ...828...83G", "2016ApJ...832...23Y", "2016ApJ...833..237M", "2016RAA....16....1H", "2016SoPh..291.1887C", "2017A&A...597A.104V", "2017A&A...604A..36G", "2017ApJ...839...67S", "2017ApJ...839..128W", "2017ApJ...840...84S", "2017ApJ...842....3N", "2017ApJ...845...18Y", "2017ApJ...851..111S", "2017ApJS..229....3C", "2017ScChD..60.1408G", "2017SoPh..292...29L", "2017SoPh..292...36L", "2017SoPh..292..105Y", "2018A&C....24..117M", "2018ApJ...855...11H", "2018ApJ...856...79Y", "2018ApJ...860..144S", "2018ApJ...869..110S", "2018JSWSC...8A..25L", "2019ASPC..526..255Y", "2019ApJ...881..101L", "2019ApJ...884..175W", "2019ApJ...886..149L", "2019ApJ...887...99C", "2019ApJS..243...36L", "2020A&A...635A..64G", "2020A&A...642A..52A", "2020AdSpR..65.2828H", "2020ApJ...893..159S", "2021A&A...648A..28A", "2021A&A...649A.129G", "2021A&A...652A..55W", "2021ApJ...911..130H", "2021ApJ...913..147L", "2021ApJ...918...15C", "2021ApJ...923...84G", "2021ApJ...923..151A", "2021ApJS..256...26B", "2021JSWSC..11....4P", "2021JSWSC..11...14P", "2021RNAAS...5..134S", "2021RSPTA.37900175N", "2022A&A...666A..93H", "2022ApJ...930...60H", "2022ApJ...936...37L", "2022ApJ...936...94S", "2022ApJ...938..117N", "2022ApJS..259...24H", "2022Mine...12.1158L", "2022NatCo..13..640Y", "2022SoPh..297...17L", "2022SoPh..297..121L", "2022tess.conf41007L", "2023A&A...669A.144S", "2023ApJ...942...27L", "2023ApJ...942...83D", "2023ApJ...944..111A", "2023ApJ...951...23C", "2023ApJ...954..199L", "2023ApJ...954L..20L", "2023ApJ...955...40L", "2023ApJS..264...49F", "2023ApJS..269...54W", "2023BAAS...55c.335R", "2023CosRe..61..265S", "2023MNRAS.524.5880Y", "2024ApJ...961..148L", "2024arXiv240615272G" ]
[ "astronomy" ]
8
[ "Helioseismic and Magnetic Imager", "Solar Dynamics Observatory", "Photosphere", "Sun", "magnetic fields", "Techniques", "polarimetric", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1956PASJ....8..108U", "1962IzKry..28..259R", "1969drea.book.....B", "1987ApJ...322..473S", "1990ApJ...348..747L", "1992nrfa.book.....P", "1997ApJ...474..810M", "2006A&A...456.1159M", "2007MmSAI..78..148L", "2007SoPh..240..177B", "2011A&A...527A..29B", "2011SoPh..273..267B", "2012SoPh..275....3P", "2012SoPh..275..207S", "2012SoPh..275..229S", "2012SoPh..275..285C", "2012SoPh..277...89L", "2012SoPh..279..295L" ]
[ "10.1007/s11207-014-0497-7", "10.48550/arXiv.1403.3677" ]
1403
1403.3677_arXiv.txt
The Very Fast Inversion of the Stokes Vector \citep[VFISV;][]{Borr11} is a spectral line inversion code tailored and optimized to invert the full disk spectro-polarimetric data of the {\it Helioseismic and Magnetic Imager} (HMI) instrument \citep{Scherrer2012, Schou2012} on board the {\it Solar Dynamics Observatory} \citep[SDO;][]{Pesnell2012}. HMI is a filtergram-type instrument that observes (with one camera) the Stokes {\it I}, {\it Q}, {\it U}, and {\it V} at six wavelength positions across the Fe {\sc i} 6173 \AA\ spectral line for the full disk of the Sun with 16 million pixels (4096$\times$4096 CCD) every 135 s. All of the data pipeline procedures, and the spectral line inversion code in particular, have limited allowable runtimes in order to keep pace with the data flow rate and prevent a processing backlog. In the forward problem, VFISV solves the radiative transfer equation (RTE) for polarized light using the Milne-Eddington (ME) approximation to generate a set of Stokes profiles from a given model atmosphere \citep[][]{unno, rachkovsky}. The ME approximation assumes that all the parameters describing the atmosphere are constant along the line of sight (LOS) except for the source function, which varies linearly with optical depth. In addition, the generation of polarized radiation is formulated in the classical Zeeman effect regime. Traditionally, ME models applied to polarized RTE problems use up to eleven free parameters to describe the atmosphere in which the Stokes profiles are generated. There are five thermodynamical parameters: the line-to-continuum absorption ratio, $\eta_0$, the Doppler width, $\Delta\lambda_{\rm D}$, the damping parameter of the Voigt function, $a$, and the components of the linearized source function, $B_0$ and $B_1$, respectively. The magnetic field vector is described by three variables, {\it i.e.} the magnetic field strength, $B$, its inclination with respect to the LOS, $\gamma$, and its azimuth in the plane perpendicular to the line of sight, $\psi$ (which is referenced to a column of pixels on the HMI CCD and increases counter-clockwise). There are also two kinematic parameters: the Doppler velocity, $v$, and the macroturbulent velocity, $v_{\rm mac}$. The former characterizes the macroscopic plasma speed, while the latter is generally used to model a combination of unresolved plasma velocity fields and instrument smearing effects - it is usually expressed in the form of a convolution of the Stokes profiles with a gaussian function of width $v_{\rm mac}$. VFISV makes explicit use of the measured HMI transmission filter profiles in the spectral line synthesis and hence does not need to use $v_{\rm mac}$ to account for the instrumental broadening, while the other thermodynamic parameters compensate for the unresolved velocities and any residual instrumental broadening. A standard additional geometrical parameter known as the filling factor, $\alpha$, quantifies the fraction of light within any given pixel that originates from a magnetized atmosphere. The optimization scheme of VFISV, based on a Levenberg-Marquardt (LM) minimization algorithm \citep[see][]{press}, takes a set of observed Stokes profiles and finds the model parameters that best describe the atmosphere in which they were generated. It achieves this by performing a non-linear minimization of a merit function, $\chi^2$, that measures the similarity between the observed and synthetic Stokes profiles in a least squares sense. VFISV is just one module among many in the data pipeline for the HMI instrument. It operates on Level1.5 data (the \verb|hmi_S.720s| data-series) from the 'vector' camera of HMI averaged every 720 s \citep[][]{hoeksema2014} and its output is fed to a disambiguation code \citep[see][]{barnes2014} that resolves the $180^\circ$ azimuth ambiguity of the magnetic field vector. The original version of VFISV \citep{Borr11} was developed before the launch of SDO, which took place on 11 February 2010. The spectral line inversion code has been further optimized since HMI data became available online. The purpose of this paper is to describe the changes implemented in the code and its output since the launch of SDO. The current version of the code is referred to as \verb|fd10|, which began as a number (in this case 10) assigned to identify tests of various versions of the code on the full-disk (FD) data. The data series produced by VFISV \verb|fd10| in the HMI pipeline processing is called \verb|hmi.ME_720s_fd10| and is available through JSOC\footnote{http://jsoc.stanford.edu/} (Joint Science Operations Center). Its name reflects the fact that it is a product of the Milne-Eddington inversion (ME) every 720 s with the \verb|fd10| version of the VFISV code. This manuscript is organized as follows. Section \ref{section:chi2-space} reports on the changes in the code that alter the $\chi^2$-space and hence determine which families of solutions are preferred over others, whilst Section \ref{section:speed} describes all of the procedures and alterations that have a direct impact on the speed performance of the inversion but not on the final solution found by the algorithm. Section \ref{section:conclusions} sums up the improvements in speed and performance of the code and some bug fixes are reported in the Appendix.
\label{section:conclusions} Some of the changes to the VFISV code reported in this paper were aimed at speeding up the code while having little or no impact on the accuracy of the solution found by the inversion algorithm. These include measures that prevent the algorithm from performing unnecessary calculations and from following inefficient paths through the parameter space in pursuit of the solution. Certain modifications tried to palliate other performance issues that arise from the degeneracy among the atmospheric model parameters (which lead to multiple minima in the $\chi^2$) from the sometimes inadequate simplicity of the model chosen to represent the complexity of the real Sun, and from the data themselves, which have limited spectral and spatial resolution, are affected by photon noise, {\it etc}. Throughout this manuscript we organized the changes into two classes: those that altered the shape of the $\chi^2$ function (and hence the model atmosphere), and those that did not. However, they are not all independent of one another, and they operate together to improve the overall efficiency and performance of VFISV in the HMI data pipeline. \begin{itemize} \item A set of weights for the Stokes profiles that provided the smoothest possible solution inside active regions was chosen. This renders less than optimal results in weakly magnetized areas where the polarization profiles are heavily dominated by the noise. Because of the smearing effect of the transmission filter profiles of the HMI instrument, the standard weighting system used to invert datasets from spectrograph type of instruments (such as ASP or {\it Hinode}/SOT-SP), does not work with HMI data. A custom set of weights was chosen and applied uniformly over the entire FOV. \item A regularization term that penalizes high values of $\eta_0$ was added to the merit function. This mitigates the double minima problem in the parameter space and helps prevent the code from converging to unphysical solutions. \item Limits on the range of each atmospheric parameter were set to prevent the algorithm from probing extreme unphysical values. \item A diagnostic procedure that identifies possibly problematic pixels enforces a series of inversion restarts that attempt to localize the global minimum. Coupled to this, a discriminating convergence criteria ensures a very accurate solution while keeping the average number of iterations relatively low ($\approx 30$). \item A major speed improvement (about 64\% reduction in computing time) was achieved, without compromising the accuracy of the solution, by considering a spectral range of $\pm 2$ \AA\ around the central wavelength but limiting the explicit forward modeling calculation to an inner $\pm 0.65$ \AA\ range. This hybrid approach accounts for the light that passes through the secondary lobes of the HMI transmission filter profiles without doing the detailed spectral line calculation far out into the continuum. \item An additional 10\% increase in speed was obtained after performing two changes of variable on the atmospheric model parameters before solving the linear system of equations at each iterative step. This modification prevented the code from performing useless iterations in the parameter space by helping it find a more direct route to the solution. \item Coding errors were identified and corrected. Some of these are reported in the Appendix. \end{itemize} \begin{acks} Much of the reported effort was supported by Stanford University NASA Grant NAS5-02139 for SDO/HMI commissioning and pipeline code implementation; KDL and GB also acknowledge PO\# NNG12PP28D/C\# GS-23F-0197P from NASA/Goddard Space Flight Center. The National Center for Atmospheric Research is sponsored by the National Science Foundation. \end{acks} \appendix % \label{section:bugs} This section reports some coding mistakes that have been spotted in VFISV since its original release and that have been fixed in the \verb|fd10| HMI data pipeline version. This version can be found through the JSOC (Joint Science Operations Center) CVS tree at Stanford.
14
3
1403.3677
The Very Fast Inversion of the Stokes Vector (VFISV) is a Milne-Eddington spectral line inversion code used to determine the magnetic and thermodynamic parameters of the solar photosphere from observations of the Stokes vector in the 6173 Å Fe I line by the Helioseismic and Magnetic Imager (HMI) onboard the Solar Dynamics Observatory (SDO). We report on the modifications made to the original VFISV inversion code in order to optimize its operation within the HMI data pipeline and provide the smoothest solution in active regions. The changes either sped up the computation or reduced the frequency with which the algorithm failed to converge to a satisfactory solution. Additionally, coding bugs which were detected and fixed in the original VFISV release are reported here.
false
[ "active regions", "SDO", "Magnetic Imager", "HMI", "the original VFISV inversion code", "Stokes", "observations", "the HMI data pipeline", "a Milne-Eddington spectral line inversion code", "the Solar Dynamics Observatory", "the smoothest solution", "order", "a satisfactory solution", "the original VFISV release", "the Stokes vector", "the Stokes Vector", "the Helioseismic and Magnetic Imager", "the magnetic and thermodynamic parameters", "(HMI", "Milne-Eddington" ]
11.972084
15.320716
2
746712
[ "Agarwal, Bhaskar", "Dalla Vecchia, Claudio", "Johnson, Jarrett L.", "Khochfar, Sadegh", "Paardekooper, Jan-Pieter" ]
2014MNRAS.443..648A
[ "The First Billion Years project: birthplaces of direct collapse black holes" ]
97
[ "Max-Planck-Institut für extraterrestrische Physik, Giessenbachstraße, D-85748 Garching, Germany", "Max-Planck-Institut für extraterrestrische Physik, Giessenbachstraße, D-85748 Garching, Germany; Instituto de Astrofísica de Canarias, C/ Vía Láctea s/n, 38205 La Laguna, Tenerife, Spain; Departamento de Astrofísica, Universidad de La Laguna, Av. del Astrofísico Franciso Sánchez s/n, 38205 La Laguna, Tenerife, Spain", "Max-Planck-Institut für extraterrestrische Physik, Giessenbachstraße, D-85748 Garching, Germany; X Theoretical Division, Los Alamos National Laboratory, Los Alamos, NM 87545, USA", "Max-Planck-Institut für extraterrestrische Physik, Giessenbachstraße, D-85748 Garching, Germany; Institute for Astronomy, University of Edinburgh, Royal Observatory, Edinburgh, EH9 3HJ, UK", "Max-Planck-Institut für extraterrestrische Physik, Giessenbachstraße, D-85748 Garching, Germany" ]
[ "2014A&A...572A..22V", "2014ApJ...792...78L", "2014ApJ...795..137R", "2014MNRAS.442.2036D", "2014MNRAS.443.1979L", "2014MNRAS.445..544S", "2014MNRAS.445..686J", "2014MNRAS.445.1056V", "2015ApJ...799...98M", "2015ApJ...805...73I", "2015ComAC...2....3G", "2015MNRAS.446.2380B", "2015MNRAS.446.3163L", "2015MNRAS.448..104P", "2015MNRAS.449.3766R", "2015MNRAS.450.4411C", "2015MNRAS.452..755S", "2015MNRAS.452.1233H", "2015MNRAS.454.3771P", "2015NatCo...6.7099G", "2016A&A...585A..11S", "2016ASSL..423...65M", "2016ApJ...817...20M", "2016ApJ...827...58S", "2016ApJ...829L...6S", "2016MNRAS.455....2F", "2016MNRAS.456.1901H", "2016MNRAS.457.2423Y", "2016MNRAS.457.3356V", "2016MNRAS.459.3377R", "2016MNRAS.459.4209A", "2016MNRAS.460.4003A", "2016MNRAS.461..111R", "2016MNRAS.462.2184H", "2016MNRAS.462.3591M", "2016MNRAS.463..529H", "2016MNRAS.463..870S", "2017A&A...603A..69G", "2017ApJ...838..111Y", "2017ApJ...838..117N", "2017IAUS..324..257N", "2017IJMPD..2630021M", "2017MNRAS.467.4293C", "2017MNRAS.468L..82A", "2017MNRAS.469..231A", "2017MNRAS.469.2517W", "2017MNRAS.469.3329W", "2017MNRAS.470.1121T", "2017MNRAS.470.4034A", "2017MNRAS.471.4878S", "2017MNRAS.472..205S", "2017MNRAS.472.2773G", "2017MNRAS.472.4414D", "2017NatAs...1E..75R", "2017PASA...34...31V", "2018A&A...614A.121N", "2018AJ....155..131M", "2018ApJ...861...39D", "2018JCAP...05..017B", "2018MNRAS.474.3825V", "2018MNRAS.475.4252A", "2018MNRAS.475.4636R", "2018MNRAS.476..407V", "2018MNRAS.477.3030K", "2018MNRAS.478.3756C", "2018MNRAS.478.5037R", "2018MNRAS.479L..23H", "2018MNRAS.480.5673E", "2018RvMP...90b5006K", "2018arXiv180306430C", "2019ApJ...874..117B", "2019ConPh..60..111S", "2019MNRAS.486.2336D", "2019MNRAS.488.2006E", "2019MNRAS.488.3268A", "2019PASA...36...27W", "2020ApJ...897..103S", "2020MNRAS.494.2851C", "2020MNRAS.497.3761L", "2021MNRAS.500..138W", "2021MNRAS.501.1413N", "2021MNRAS.501.2029C", "2021MNRAS.508.1756L", "2021MNRAS.508.3226A", "2021arXiv210313932N", "2022MNRAS.510..177B", "2022MNRAS.514.4912H", "2022NewAR..9401642M", "2022PASP..134l1001M", "2023ApJ...952..146N", "2023LRR....26....2A", "2023MNRAS.518.4606P", "2023MNRAS.518.4672S", "2023MNRAS.520.2081P", "2023MNRAS.522..330I", "2023MNRAS.525.4832Y", "2024ApJ...960L...1N" ]
[ "astronomy" ]
6
[ "galaxies: high-redshift", "quasars: supermassive black holes", "dark ages", "reionization", "first stars", "early Universe", "Astrophysics - Cosmology and Nongalactic Astrophysics", "Astrophysics - Astrophysics of Galaxies" ]
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[ "10.1093/mnras/stu1112", "10.48550/arXiv.1403.5267" ]
1403
1403.5267_arXiv.txt
The observation of multiple high redshift quasars at $z>6$ \citep[e.g.][]{Mortlock:2011p447,Venemans:2013p3261} demands an explanation for the origin and growth of the (supermassive) black holes (BHs) fuelling quasar activity. The idea of the direct collapse of pristine gas in primordial haloes into BHs, i.e. direct-collapse black holes (DCBHs), is aimed at solving this problem by providing a physical mechanism to form seed BHs with masses $\Mbh \sim 10^{4-6}\msun$ \citep[e.g.][]{Eisenstein:1995p870}. The main advantage of the DCBH scenario is a much larger seed mass than what is expected from population III (hereafter, Pop~III) stellar remnants ($\sim 100\msun$), which makes it easier to grow to supermassive scales in a relatively short time \citep[e.g.][]{Volonteri:2008p1043}. During collapse the gas cloud must avoid fragmentation and lose its angular momentum in order to form a high density gas core. The core could ultimately results in a DCBH if it can accrete at a rate of $\sim 0.1$-$1\msunyr$ for $10^{5-6}\Myr$ \citep[e.g.][]{Latif:2013p2446}. That said, a pristine, low spin, atomic cooling halo with a critically low H$_2$ fraction is the prerequisite for DCBH formation \citep{Bromm:2003p22}. Pristine gas is necessary as the injection of metals in the halo (e.g.~from neighbouring supernova driven winds) would lead to fragmentation (and eventually star formation) during the collapse process, as metals are effective coolants. A low spin halo can allow for the efficient transport of angular momentum, thereby facilitating the accumulation of gas towards the central region \citep{Lodato:2006p375}. A low H$_2$ fraction ensures that the pristine gas cools mostly \textit{via} atomic hydrogen, which sets the minimum cooling temperature at $\sim 8000\K$, thereby raising the Jeans mass required for collapse to $10^5\msun$ at a density of $10^4\ \rm cm^{-3}$. The suppression of H$_2$ cooling can be attained through a high level of Lyman-Werner (LW) radiation as it can effectively dissociate H$_2$ into atomic hydrogen \citep{Shang:2010p33,WolcottGreen:2011p121}. Following the prerequisites outlined above, the DCBH can form through various channels: the supermassive star stage \citep{Begelman:1978p792,Begelman:2010p872}, the quasi-star stage \citep[e.g.][]{Begelman:2008p672}, or \textit{via} runaway collapse from a gas disc \citep{Koushiappas:2004p871,Lodato:2006p375}. To understand DCBH formation, one must first probe its plausibility, i.e.~the conditions required for a halo to qualify as a direct-collapse (DC) candidate. \cite{Dijkstra:2008p45} (D08 hereafter) used Monte Carlo merger trees to predict the existence of such sites in the high-redshift Universe. They employed two-point correlation functions and halo mass functions and predicted a few DC sites per co-moving Gpc$^3$ volume. A recent study by \cite{Agarwal:2012p2110} (A12 hereafter) used a suite of semi-analytical models, on top of a cosmological N-body simulation, to predict the abundance of DC sites at $z\sim6$. Their model included tracking halo histories using merger trees and the spatial variation of LW radiation from both Pop~III and Pop~II stars. They predicted as many as few DC sites per co-moving Mpc$^3$, which is significantly higher than the earlier estimate (D08) and mainly due to taking into account the local variation in LW flux due to clustered star formation and the revision in the value of the critical level of LW radiation required to cause DC \citep{Shang:2010p33,WolcottGreen:2011p121}. Although A12 was an improvement over the earlier estimates of abundances of DC sites, the model was missing metal dispersion in the inter-galactic medium due to supernova driven winds, the self consistent treatment of gas physics (e.g. cooling, dissociation, photoionisation) and thereby could have been an overestimate. Several hydrodynamical simulations have been employed to study the processes by which gas lose angular momentum and lead to the formation of a dense cloud that can undergo runaway gravitational collapse \citep[e.g.][]{Oh:2002p836,Bromm:2003p22,Begelman:2006p75}. Turbulence has been found to be one of the main agents via which gas can accumulate at the centre of metal-free atomic cooling haloes \citep{Wise:2008p2441,Latif:2013p2446}. However, the formation of a galactic-type disc has also been reported \citep{Regan:2009p776}. Note that, in all these studies, single, isolated haloes, extracted from cosmological simulations in some cases, were assumed \textit{a priori} to fullfill the criteria for DC. The strength of the current study is the use of a hydrodynamical, cosmological simulation that self-consistently accounts for pair-instability (PISN) and core-collapse SN feedback from Pop~III and Pop~II stars (in the form of enrichment and energetic feedback), the self-consistent evolution of the global and local photo-dissociating LW radiation from both stellar populations, and the photo-ionisation of atomic ($\mathrm{H}^-+\gamma\rightarrow \mathrm{H} + \mathrm{e}$) and photo-dissociation of molecular ($\mathrm{H}_2+\gamma\rightarrow \mathrm{H} + \mathrm{H}$) hydrogen species. The advantage of such approach is that, for the selected candidate haloes, we know their formation history and the environment they live in. The paper is organised as follows. We briefly describe the FiBY simulation used in this work and the modelling of LW radiation and self-shielding in section~\ref{sec.method}. The results of our study are presented in section~\ref{sec.results} where we discuss the nature of the DC sites, their merger histories and the nature of the galaxies in their local neighbourhood. The summary of the work and the discussion of the results are presented in section~\ref{sec.summary}.
\label{sec.summary} In this study, we have employed one of the FiBY project's simulations to pin-point the location and environment of metal-free, non star forming, atomic cooling haloes within a cosmological hydrodynamical simulation. We report a LW flux that is considerably higher than the global mean in 6 candidates, which makes them potential candidates for DCBH formation sites. Our attempt was to quantify if such sites could exist in a cosmological simulation that forms Pop III and Pop II stars self-consistently, accounts for a self--consistent build up of local and global LW radiation flux, and includes metal dispersion via SNe and stellar winds. In order to identify such DC sites, we first identified a sample of pristine, atomic cooling, non star forming haloes and then selected the ones that are exposed to the highest levels of LW radiation, as outlined in Sec. \ref{sec.dc selection}. The sample of 6 haloes identified in this study hints towards the haloes being possible sites of DCBH, however, the formation of a DCBH would depend on the state of the subsequent gas collapse, which could be probed by extracting these haloes and simulating them in a zoom hydrodynamical simulation that has a high enough resolution. The fact that we have a handful of potential DC sites in our 4 cMpc side-box suggest that DCBHs do not need high-$\sigma$ regions to form, and in fact, can even form in stand alone haloes that happen to be in the vicinity of a few modestly star forming galaxies that cumulatively produce the critical level of LW radiation (see DC3, Fig.\ref{fig.mergertrees}). We report that satellite haloes of larger galaxies are the most likely sites for DCBH formation, i.e. massive seed BHs form outside the galaxies they eventually end up in (A12, A13). The critical value of the LW specific intensity that favours the formation of DCBHs has been derived in the literature by studying haloes in isolation, where an ideal atomic-cooling halo is selected from a cosmological setup and is irradiated with an increasing level of LW flux till a point is reached where H$_2$ cooling becomes insufficient \citep{Bromm:2003p22,Shang:2010p33,WolcottGreen:2011p121}. This is an assumption on the physical conditions, as the halo under question would be subject to time varying LW feedback from neighbouring galaxies right since its birth. Therefore the previous calculations of $J_{\rm crit}$ might have been overestimated. The exposure of the halo to a time--varying LW flux, ever since its birth, was self-consistently accounted for in the FiBY simulation analysed in this study. \footnote{However note that the authors used reaction rates based on a black body spectrum with a temperature of $10^4$ and $10^5\ \rm K$ representative of a Pop II and Pop III stellar population respectively.} Therefore whether or not the rest of the haloes (besides the 6 DC candidates) in the sample that are exposed to lower values of $J_{\rm LW}$ could harbour a DCBH is unclear, as the only criterion the other haloes do not meet is the exposure to high levels of LW radiation. The subsequent accretion process and the final mass that these DCBHs attain would be highly dependent on the mergers that the DC haloes go through. As mentioned earlier, DC0, DC1, DC4, DC5 form as satellites of a larger galaxies and eventually undergo mergers. Upon formation, the DCBH could engulf a major fraction of the gas in its host galaxy \citep[see for e.g.][]{Schleicher:2013p2775}, thereby running out of gas for subsequent accretion. Mergers with larger gas rich galaxies could turn on the accretion process again, aiding these DCBHs to attain supermassive scales (A13). Surprisingly, DC3 does not form as a satellite of a larger galaxy and the host halo attains a mass of $\approx 10^{7.5}\msun$ by the end of the simulation ($z=6$). Whether this particular candidate evolves into a Milky Way type galaxy, or ends up in a scenario of quenched DCBH accretion due to insufficient fuelling is unclear and the subject of an undergoing study. Note that DC4 and DC5 end up in the same galaxy at $z=6$, hinting towards the possible event of a DCBH merger in the early Universe. This sort of event, would be an ideal candidate in explaining the growth of massive seed BHs to supermassive scales, where upon undergoing a merger, the seeds could double their mass and continue to grow by rapid gas accretion. However, mergers are subject to gravitational recoil and dynamical friction effects, which could hinder the growth of these DCBHs. In order to find such an optimised event, one would need to run our simulation with a much larger box size, a feat unattainable with the current computational limitations. That said, the occurrence of these sites in our relatively-small simulation volume hints towards the possibility that most present day galaxies might be harbouring a DCBH at their centres. In this study we haven't touched upon the effects of reionisation on the DCBH formation sites. The candidate haloes cross the atomic cooling limit at redshifts between 8 and 10, when the CMB data suggests that reionisation was already under way \citep{Komatsu:2011p409,Ade:2013p2565}. Moreover, the local sources that produce the high level of LW radiation at the potential formation sites likely also produce a high level of ionising radiation. Thus, it is possible that reionisation affects the candidate DCBH sites. The effect of reionisation on DCBH formation is unclear. Ionising photons could heat the gas in haloes that are unable to shield from the radiation, preventing the haloes from growing and inhibiting collapse to a DCBH. Also, the additional free electrons allow for faster H$_2$ formation through the H$^-$ channel at low densities, causing the gas inside the halo to cool more efficiently. Since the effect of reionisation on DCBH formation depends critically on the ability of the formation sites to self-shield against the ionising radiation, addressing this issue requires simulations with a higher resolution than we employ here. We have therefore decided to address these issues in a follow-up study (Johnson et al in prep). The next step is to extract these haloes and simulate them for their entire formation histories, with the associated LW radiation (and other properties) as input, in high-resolution zoom simulations. This will shed new light on the role that LW radiation and an ionising flux, amongst other properties, plays on the process of DCBH formation.
14
3
1403.5267
We investigate the environment in which direct-collapse black holes may form by analysing a cosmological, hydrodynamical simulation that is part of the First Billion Years project. This simulation includes the most relevant physical processes leading to direct collapse of haloes, most importantly, molecular hydrogen depletion by dissociation of H<SUB>2</SUB> and H<SUP>-</SUP> from the evolving Lyman-Werner radiation field. We selected a sample of pristine atomic-cooling haloes that have never formed stars in their past, have not been polluted with heavy elements and are cooling predominantly via atomic hydrogen lines. Amongst them we identified six haloes that could potentially harbour massive seed black holes formed via direct collapse (with masses in the range of 10<SUP>4-6</SUP> M<SUB>⊙</SUB>). These potential hosts of direct-collapse black holes form as satellites are found within 15 physical kpc of protogalaxies, with stellar masses in the range ≈10<SUP>5-7</SUP> M<SUB>⊙</SUB> and maximal star formation rates of ≈0.1 M<SUB>⊙</SUB> yr<SUP>-1</SUP> over the past 5 Myr, and are exposed to the highest flux of Lyman-Werner radiation emitted from the neighbouring galaxies. It is the proximity to these protogalaxies that differentiates these haloes from rest of the sample.
false
[ "maximal star formation rates", "stars", "atomic hydrogen lines", "direct collapse", "≈0.1 M<SUB>⊙</SUB", "massive seed black holes", "heavy elements", "haloes", "stellar masses", "direct-collapse black holes form", "direct-collapse black holes", "pristine atomic-cooling haloes", "Lyman-Werner radiation", "masses", "the First Billion Years project", "the evolving Lyman-Werner radiation field", "dissociation", "the neighbouring galaxies", "protogalaxies", "Lyman-Werner" ]
13.998363
6.917088
119
745827
[ "Rowlands, K.", "Gomez, H. L.", "Dunne, L.", "Aragón-Salamanca, A.", "Dye, S.", "Maddox, S.", "da Cunha, E.", "van der Werf, P." ]
2014MNRAS.441.1040R
[ "The dust budget crisis in high-redshift submillimetre galaxies" ]
104
[ "(SUPA) School of Physics and Astronomy, University of St Andrews, North Haugh, St Andrews, Fife KY16 9SS, UK; School of Physics and Astronomy, The University of Nottingham, University Park Campus, Nottingham NG7 2RD, UK", "School of Physics and Astronomy, Cardiff University, Queens Buildings, The Parade, Cardiff CF24 3AA, UK", "Department of Physics and Astronomy, University of Canterbury, Private Bag 4800, Christchurch 8140, New Zealand", "School of Physics and Astronomy, The University of Nottingham, University Park Campus, Nottingham NG7 2RD, UK", "School of Physics and Astronomy, The University of Nottingham, University Park Campus, Nottingham NG7 2RD, UK", "Department of Physics and Astronomy, University of Canterbury, Private Bag 4800, Christchurch 8140, New Zealand", "Max Planck Institute for Astronomy, Konigstuhl 17, D-69117 Heidelberg, Germany", "Leiden Observatory, Leiden University, P.O. Box 9513, NL-2300 RA Leiden, the Netherlands" ]
[ "2014A&A...568A..92M", "2014A&A...569A..98T", "2014MNRAS.441.1017R", "2014MNRAS.444.2442V", "2014Natur.511..296G", "2015A&A...577A..80M", "2015A&A...582A.121R", "2015ApJ...799...50L", "2015ApJ...800...50M", "2015ApJ...813..119D", "2015ApJ...814...17M", "2015EAS....71...57G", "2015MNRAS.446.1512H", "2015MNRAS.447.2937H", "2015MNRAS.448..258R", "2015MNRAS.449.1625B", "2015MNRAS.451.2030D", "2015MNRAS.452..397C", "2016A&A...587A.157B", "2016ASSL..423...65M", "2016ApJ...818...62S", "2016MNRAS.457.3775M", "2016MNRAS.463.1291I", "2016pas..conf..179M", "2017A&A...597A...5M", "2017A&A...603A...4S", "2017ApJ...851..152B", "2017IAUS..331..290C", "2017MNRAS.464.4680D", "2017MNRAS.465...54C", "2017MNRAS.465.3309D", "2017MNRAS.471.1743D", "2017MNRAS.471.4615G", "2017MNRAS.472.2289F", "2017MmSAI..88..420V", "2017PASA...34...31V", "2018A&A...620A.112B", "2018ApJ...868...62G", "2018MNRAS.473.4538G", "2018MNRAS.475.2891D", "2018MNRAS.476.1705S", "2018MNRAS.476.3218C", "2018MNRAS.479..536D", "2018MNRAS.479.1077B", "2018MNRAS.479.2588G", "2018MNRAS.480..108Z", "2018PASA...35....2V", "2019A&A...623A...5D", "2019A&A...632A..43M", "2019ApJ...871..151R", "2019ApJ...886...51C", "2019ApJS..244...40L", "2019BAAS...51c.378C", "2019IAUS..343..406G", "2019MNRAS.483...70C", "2019MNRAS.485..440P", "2019MNRAS.488..978J", "2019MNRAS.489.4389L", "2019MNRAS.489.5256C", "2019MNRAS.489L.108D", "2019MNRAS.490.1425L", "2020A&A...641A.168N", "2020ApJ...889..137M", "2020ApJ...890...19S", "2020ApJ...903..148L", "2020IAUS..341..196F", "2020MNRAS.491.4334M", "2020MNRAS.493.2706C", "2020MNRAS.494.2355P", "2020MNRAS.496.3668D", "2020MNRAS.498.5652K", "2020MNRAS.499.4666F", "2021A&A...648A..52L", "2021A&A...649A..18G", "2021A&A...652A..23F", "2021AJ....162...83C", "2021ApJ...907...29C", "2021ApJ...910...95R", "2021MNRAS.504.4143S", "2021MNRAS.505.3228D", "2021RNAAS...5..288M", "2022ApJ...928...90R", "2022HabT.........1G", "2022MNRAS.510.3734D", "2022MNRAS.515.5306R", "2022MNRAS.517.5930S", "2023A&A...678A..28B", "2023ApJ...944...33V", "2023ApJ...946...42C", "2023IJAsB..22..347T", "2023MNRAS.518.3211S", "2023MNRAS.518.5522C", "2023MNRAS.522.2279R", "2023MNRAS.522.4852P", "2023MNRAS.525.1839F", "2023arXiv230507743P", "2024A&A...683A.232B", "2024A&ARv..32....2S", "2024ApJ...964..129M", "2024MNRAS.527.8755V", "2024NatCo..15.1841K", "2024arXiv240301202L", "2024arXiv240319798E", "2024arXiv240410568L" ]
[ "astronomy" ]
11
[ "dust", "extinction", "ISM: evolution", "galaxies: evolution", "galaxies: high redshift", "submillimetre: galaxies", "Astrophysics - Astrophysics of Galaxies", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
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[ "10.1093/mnras/stu605", "10.48550/arXiv.1403.2995" ]
1403
1403.2995_arXiv.txt
\label{Intro} The first blind submillimetre surveys discovered a population of highly star-forming ($100-1000\,$M$_\odot$yr$^{-1}$), dusty galaxies at high redshift (\citealp*{Smail97}; \citealp{Hughes98, Barger98, Eales99}). These submillimetre galaxies (SMGs) are thought to be undergoing intense, obscured starbursts \citep{Greve05, Alexander05, Tacconi06, Pope08}, which may be driven by gas-rich major mergers \citep[e.g.][]{Tacconi08, Engel10, Wang11, Riechers11}, or streams of cold gas \citep{Dekel09, Dave10, vandeVoort11a}. Observational studies show that SMGs typically have stellar masses of $\sim 10^{11}\,\rm M_{\odot}$ \citep[e.g.][]{Hainline11, Magnelli12}, large dust masses ($\sim 10^{8-9}\,\rm M_{\odot}$; \citealt{Santini10, Magdis12, Simpson14}), high gas fractions ($30-50$\,per\,cent; \citealt{Tacconi08,Bothwell12}) and Solar or sub-Solar metallicities \citep{Swinbank04, Banerji12, Nagao12}. The source of interstellar dust in SMGs is still a controversial issue, particularly whether it originates from supernovae (SNe) or from the cool, stellar winds of low-intermediate mass stars (LIMS). Recent work has revealed a \emph{`dust budget crisis'} (\citealp[][hereafter ME03]{ME03}; \citealp*{Dwek07}; \citealp*{MWH10}; \citealp{Michalowski10b, Santini10}; \citealp*{Gall11a}; \citealp{Valiante11}), whereby it is difficult to explain the high dust masses observed in high redshift galaxies through dust from LIMS\footnote{Note that this problem is not limited to high redshift SMGs, and is also seen in galaxies at low redshift \citep[e.g.][]{Matsuura09, Dunne11, Rowlands12, Smith_HRS12, Boyer12}.}. At $z>5$ this is further compounded as there is little time for LIMS to produce significant amounts of dust (ME03; \citealt{Dicris13}). The surprisingly constant dust-to-metals ratio measured in galaxies over a wide range of cosmic time also indicates that a rapid mechanism of dust formation is needed \citep*[][and references therein]{Zafar13}, requiring dust formation timescales to be the same order as the metal enrichment timescale. Although \citet{Valiante09} and \citet{Dwek11} argue that AGB stars may contribute significantly to the dust budget after only $150-500$~Myrs (and thus may be a significant source of dust at high redshift), the amount of dust produced is highly sensitive to the assumed initial mass function (IMF). Furthermore, in the former study a high star-formation rate (SFR) in excess of $1000\,$M$_\odot$yr$^{-1}$ sustained over $\sim 0.3-0.4$\,Gyr is required to build up a significant mass of dust. Due to their short lifetimes, massive-star SNe have long been proposed as a potential source of dust at early times (ME03; \citealp{Nozawa03, Dunne03a, Dwek07, Dunne09b}; \citealp*{Gall11Rev}). Observational evidence of dust formation in SN ejecta has come to light recently with SN1987A, Cas A and the Crab Nebula remnants containing significant quantities of dust ($0.1-1\,\rm M_{\odot}$; \citealp{Dunne09b,Matsuura11,Gomez12b}). There is now little doubt that dust is formed in SN ejecta \citep{Dunne03a,Sugerman06, Rho08, Barlow10, Matsuura11, Temim12, Gomez12b} though the amount of dust which will ultimately survive the supernova shocks is still highly uncertain \citep{Bianchi07, Kozasa09, Jones_Nuth11}. Additionally, dust grain growth in the ISM \citep{Draine09} has been proposed as an extra source of dust in galaxies at both high redshift \citep{MWH10,Hirashita_Kuo11,Valiante11,Calura14}, and at low redshift \citep{Dwek07, Dunne11, Inoue12, Kuo_Hirashita12, Mattsson_Andersen12b, Boyer12, Asano13}, which could make up the shortfall in the dust budget of galaxies. The difficulty with determining the origin of dust in galaxies and its lifecycle arises due to a lack of large samples of sources in which to test these issues. Previous authors including ME03; \citet{Valiante09,Valiante11,Dwek11,Gall11a} and \citet{MWH10} investigated the origin and evolution of dust in high redshift galaxies, but these were limited to one or two (or, at most, a handful) extreme starbursting systems, selected in a non-uniform way and often missing crucial far-infrared (FIR) photometry spanning the peak of the dust emission. In \citet{Rowlands14b} a sample of SMGs were carefully selected from the comprehensive data in \citet{Magnelli12} and galaxy properties were derived for the population by fitting their spectral energy distributions (SEDs) from the UV to the submillimetre in a consistent way. Here, we investigate the origin of dust in these high redshift SMGs using an updated version of the chemical evolution model of ME03 which incorporates realistic star-formation histories (SFHs) for each galaxy, with greater complexity than previous chemical evolution studies have attempted. The sample properties and derivation of the observational parameters are described in full in \citet{Rowlands14b} (see also \citealt{Magnelli12}, hereafter M12). We briefly comment on our sample selection and the spectral energy distribution (SED) fitting method in Section~\ref{sec:sample_selection}. In Section~\ref{sec:chem_ev_description} we present the updated chemical evolution model which follows the build-up of dust over time, with comparison to the observed properties of SMGs in Section \ref{sec:chem_ev}. Our conclusions are summarised in Section \ref{sec:conclusions}. We adopt a cosmology with $\Omega_m=0.27,\,\Omega_{\Lambda}=0.73$ and $H_o=71\, \rm{km\,s^{-1}\,Mpc^{-1}}$.
\label{sec:conclusions} In this paper we have used an updated chemical evolution model to reproduce the properties of a submillimetre selected sample of 26 massive, dusty galaxies in the redshift range $1.0<z<5.3$. Our chemical modelling for the first time utilises complex SFHs derived from SED fitting to the the UV--submillimetre photometry and a detailed treatment of the dust sources and sinks in galaxies. We can rule out a number of models (Table~\ref{tab:chem_ev_summary_all}) which result in dust-to-stellar masses and/or gas-to-dust ratios which are inconsistent with observations of SMGs. These models include those with dust produced by LIMS only, and those which have efficient dust destruction (mass of ISM $\mISM=1000\,$M$_\odot$ cleared of dust). The models which best match the observed gas-to-dust ratios include rapid dust build-up from grain growth and supernova dust sources. Our main results are as follows: \begin{itemize} \item We find that dust produced only by low--intermediate mass stars (LIMS) falls a factor 240 short of the observed dust masses of SMGs. Adding an extra source of dust from supernovae can account for the dust mass in SMGs in only 19\,per\,cent\, of cases. Even after accounting for dust produced by supernovae, the remaining deficit in the dust mass budget suggests that higher supernova metal yields, and/or substantial grain growth are required in order for the dust mass predicted by the chemical evolution models to match observations of SMGs. \item Efficient destruction of dust grains by supernova shocks ($\mISM=1000\,$M$_\odot$) on average decreases the dust mass from LIMS+SNe by a factor of 6--10. Additional sources of dust are required in order to account for the additional shortfall of dust in SMGs caused by dust destruction. Alternatively, dust destruction may be less efficient if dust grains are shielded from supernova shocks in dense regions of the ISM. A small amount of dust destruction ($\mISM=100\,$M$_\odot$) can be accommodated in our models only if dust is produced efficiently by \emph{both} stellar and interstellar sources. \item The average metallicity in the closed box model reaches $0.9$\,Z$_\odot$, which is consistent with the metallicity measured in SMGs. If inflows of pristine gas occur with a rate equal to the SFR the metallicity is reduced to $0.7$\,Z$_\odot$; a similar metallicity is reached with enriched gas outflows. Inflows and outflows result in a modest decrease of a factor $<1.5$ in the dust mass of SMGs. Given the current large range in observed gas-phase metallicities in SMGs, and uncertainties due to possible AGN contamination of the emission lines, we cannot currently distinguish between different inflow and outflow rates. Measurements of gas-phase metallicities which are not affected by the presence of an AGN are required for larger samples of SMGs. \item A top-heavy IMF cannot account for the observed dust masses if dust is produced by LIMS only. With no dust destruction we found that a top-heavy IMF with dust produced by both LIMS and supernovae can produce the average dust mass observed in SMGs (within a factor of two) therefore resolving the dust budget crisis. Yet, given the uncertainties involved (e.g. in the dust destruction rate and metallicity in SMGs) this does not provide unequivocal evidence for a top-heavy IMF in dusty high redshift galaxies. \item Increasing the dust emissivity on average by a factor of 7 can solve the dust budget crisis if dust is produced by LIMS and SNe and is not destroyed by supernova shocks. Variations in the dust emissivity are theoretically predicted to be a factor of $<3$, and, currently there is no observational evidence to suggest a large variation in emissivity occurs in high-redshift SMGs. Finally, an alternative explanation for the dust budget crisis is that the metal yields of stars may be systematically underestimated. \end{itemize}
14
3
1403.2995
We apply a chemical evolution model to investigate the sources and evolution of dust in a sample of 26 high-redshift (z &gt; 1) submillimetre galaxies (SMGs) from the literature, with complete photometry from ultraviolet to the submillimetre. We show that dust produced only by low-intermediate-mass stars falls a factor 240 short of the observed dust masses of SMGs, the well-known `dust-budget crisis'. Adding an extra source of dust from supernovae can account for the dust mass in 19 per cent of the SMG sample. Even after accounting for dust produced by supernovae the remaining deficit in the dust mass budget provides support for higher supernova yields, substantial grain growth in the interstellar medium or a top-heavy IMF. Including efficient destruction of dust by supernova shocks increases the tension between our model and observed SMG dust masses. The models which best reproduce the physical properties of SMGs have a rapid build-up of dust from both stellar and interstellar sources and minimal dust destruction. Alternatively, invoking a top-heavy IMF or significant changes in the dust grain properties can solve the dust budget crisis only if dust is produced by both low-mass stars and supernovae and is not efficiently destroyed by supernova shocks.
false
[ "dust", "observed SMG dust masses", "minimal dust destruction", "complete photometry", "the dust mass budget", "the dust budget crisis", "the observed dust masses", "the dust grain properties", "ultraviolet", "supernova shocks", "SMGs", "the dust mass", "substantial grain growth", "higher supernova yields", "significant changes", "IMF", "the submillimetre", "gt", "efficient destruction", "1) submillimetre galaxies" ]
12.567083
7.967348
-1
523594
[ "Yao, Lihong" ]
2014arXiv1403.0168Y
[ "A New Hypothesis On The Origin and Formation of The Solar And Extrasolar Planetary Systems" ]
0
[ "-" ]
null
[ "astronomy" ]
1
[ "Astrophysics - Earth and Planetary Astrophysics" ]
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[ "10.48550/arXiv.1403.0168" ]
1403
1403.0168_arXiv.txt
\label{newhypo} The lifetime of giant molecular clouds and massive stars may be short, but the impact of their birth and death on the surrounding intercloud (ICM) and interstellar mediums (ISM) is profound \citep[and reference therein]{yao09, zwa10}. The strong stellar winds and supernova (SN) explosions from hundreds to thousands of the massive stars create a rapid expanding hot bubble. The kinetic energy in such supersonic expansion is thermalized by a stand-off shock. The high pressure downstream drives a strong shock into the ambient ISM, or even breaks through the cloud into the adjacent ICM. These strong and fast shock waves contribute significantly to the formation and rotation of the GMCs. On the other hand, thermal and dynamical instabilities produce strong cooling and turbulence in shocked gas on small scales, which play a dominant role in providing a non-linear density perturbation through the cloud \citep[e.g.][]{hei08}. This mechanism is needed during a GMC formation to create supersonic turbulence and to allow rapid fragmentation and formation of dense clumps locally and simultaneously before the onset of global gravitational collapse of the cloud. As observed in recent studies, for example, the interaction between W44 supernova remnant (SNR) and the adjacent GMC \citep[e.g.][and references therein]{cne95, fie11, sas13}, a secondary shock wave may produce further cascade fragmentations, hence formation of smaller and higher density clumps in the cloud (see Fig.~\ref{sngmc} and Fig.~\ref{gmcsc}). These small dense cores in turn quickly collapse to form stars or a star cluster \citep{shu87, smi09}. When a cloud (or a clump) has an initial non-zero angular momentum it will collapse to an accretion disk surrounding its central mass. Observations show that the mass spectrum of the star clusters follows the same power law distribution as that of the GMCs in the ISM \citep[e.g.][]{sol97, san99, ket05}. \begin{figure} \epsscale{0.7} \plotone{f1.pdf} \caption[] {Schematic view of W44 shock wave propagates through the GMC and dense clumps detected by HCO+ J = 1 $\rightarrow$ 0 and CO J = 3 $\rightarrow$ 2 observations. Figure is obtained from Sashida et al. (2013), see details in Figure 8 of their paper. \label{sngmc}} \end{figure} \begin{figure} \epsscale{0.7} \plotone{f2.pdf} \caption[] {Schematic of rotation of a small group of dense clumps in a GMC; clump fragmentation and rotation of protostellar and protoplanet clumps. \label{gmcsc}} \end{figure} Core accretion and disk instability are the two currently accepted theories explaining how solar planets form \citep[e.g.][and reference therein]{hay85, mat07, hel13}. Core accretion model requires a few to a hundred million years to form a gas giant planet through the accretion of planetesimals from the protoplanetary disk, depending on the mass of the planet and how far it is to the parent star. This is a major problem for core accretion theory, because the time it takes to form a super giant planet (greater than a few times of the Jupiter mass) is much longer than the observed disk dissipation timescale of younger stars. For example, how can extrasolar planet GJ 504b, which has four times the Jupiter mass and about 160 million years old, form through planetesimals' collision and then core accretion at an orbital distance of 43 AU to its parent star GJ 504 \citep{kuz13}? The core accretion model also presents a problem for various planet migration theories \citep[e.g.][and reference therein]{anm11}. On the other hand, disk instability can form a super giant planet in a few hundred years, but the theory assumes that the protoplanetary disk is massive enough and can cool quickly enough to allow fragmentation to occur \citep{bos13}. In addition, sources and mechanisms for rapid fragmentation on small scales is needed to overcome global gravitational instability in the disk \citep{hei08}. Neither core accretion nor disk instability can explain the spin directions of solar planets, the conservation of angular momentum of the protostars, the evaporation of the protodisk, the formation of gas giant planet, and incredibly stable near-circular planetary orbits for billions of years, as well as why some low mass stars have planets while others have nothing around them, not even dust rings. In this study, we propose that 4.6 billion years ago our Sun along with its planets and satellites were born from a small group of dense clumps, located in one of the star clusters with thousands of stars \citep[and references therein]{cam62, elm06, zwa09}. Like many other group of dense clumps, this group of clumps was initially compressed by a strong shock wave. The densest and largest clump quickly collapsed and formed a protostar of the Sun. Other compressed clumps that orbit around the protostar will eventually collapse due to gravitational instability (i.e. disk instability theory). However, if there is a secondary shock wave, e.g. an expanding thin shell of cold gas (produced by a nearby star formation) propagates into the disk before the onset of global gravitational collapse of dense clumps in the protoplanetary disk, the dense clumps behind the shock will then be further compressed. These compressed clumps may undergo rapid cascade fragmentation which break into smaller ones, and in turn become the seeds of multiple planets and their satellites. The secondary shock wave (SW$_2$) may be weaker than the supersonic shock wave that creates the GMC and star clusters, and it should have arrived at the solar protoplanetary disk before the Sun is completely formed. This shock wave sweeps up most of the interclump gas in the protoplanetary disk into its shell before it collides with the ionized gas shock (SW$_i$) front near the radius of $R_i$ = 0.277 AU corresponding to a 80 km s$^{-1}$ of critical velocity of ionized gas particles around the forming Sun \citep{anw62, alf62}. The denser fragments behind the shock waves quickly collapse and form protoplanets and protosatellites that orbit around them, as a result of strong thermal and dynamical instabilities that dominate on the small scales \citep{hei08, bnk13}. The shock waves are partially refracted and partially reflected. Hence, part of the cold gas in the shell of SW$_2$ merges with the warm and (partially) ionized gas of SW$_i$ leaving behind the reflected shock waves an expanding non-uniform refracted gas region (0.3 AU $\rightarrow$ 10 AU), where the cooling effect converts more warm gas into the cold gas (see Fig.~\ref{shockgas}). The reflected shock waves of SW$_2$ and SW$_i$ are rarefaction waves indicated by SW$_2$$^\prime$ and SW$_i$$^\prime$ in Fig.~\ref{shockgas}. The inward expansion of the refracted gas is stalled by the high pressure plasma gas front at R$_i$, so is the reflected shock wave of SW$_i$. The outward expansion that is confined by the reflected shock wave of SW$_2$ becomes the gas supply for building gas giant planets. The dense clumps in warm gas zone (WGZ, 0.3 AU - 2.7 AU) collapse instantly and form the rocky planets, i.e. Mercury, Venus, Earth, and Mars, while collapsing clumps in neutral cold gas zone (CGZ, 2.7 AU - 10 AU) accrete sufficient cold gas to form gas giant planets Jupiter, Saturn, and their satellites. During the shock collision of SW$_2$ and SW$_i$ very small clumps (smaller than the Earth moon) that are supposed to form the satellites of rocky planets may attain or lose their orbital velocities and escape. Some of these clumps may be thrown inward toward the star where they will be heated and ionized, while other clumps may be ripped into very tiny fragments and pushed outward into the discontinuous gas region (at $\sim$ 2.7 AU) between WGZ and CGZ. Eventually these tiny fragments from WGZ along with the condensed tiny cold gas clumps form the ring of the Asteroid belt (2 AU and 4 AU). The reflected shock wave of SW$_2$ continues propagating outward and beyond Uranus, Neptune, and Pluto, then gradually slows down and eventually stalls. It soon becomes fragmented and dispersed to be part of the cold icy gas rings (i.e. Kuiper belt), when it reaches the outer region of the solar system (30 - 50 AU from the Sun). This explains why materials from the early formation of the solar system are found in the Kuiper belt, when examining stardust samples return from Comet Wild 2 \citep[and reference therein]{das11}. In addition, during the stellar core accretion, much of the angular momentum is transferred from the forming star to its surrounding ionized gas, and then from the ionized gas to the angular momentum of protoplanets and their satellites during the head-on shock wave collision. The orbital parameters (especially their semi-major axis) for all planets and their satellites have changed very little since the time of their formation. A majority of the planets and their satellites have the same spin direction as the their orbital direction and spin direction of the Sun. If these planets and their satellites, the asteroids, were formed from merely random core accretions (i.e. planetesimal accretion theory), we would see an even mixture of directions of orbital revolution and rotation of planets and their satellites. However, if there is a certain mechanism as proposed in this study, which allows a small group of dense clumps that are orbiting around the forming Sun to undergo rapid fragmentation followed by a quick collapsing both locally and simultaneously, we are expected to see the formation of multiple planets and satellites like our solar system. Hence, the newly formed planets may experience a strong spin-spin coupling and begin to align each other in the external magnetic field generated by the plasma gas surrounding the forming parent star (see Fig.~\ref{pspin}). The spin axes of the planets will precess around the magnetic field. Magnetic spin resonance is the origin of the observed rotational directions of solar planets which follows the Pascal's rule of a quintet configuration (1:4:6:4:1), as indicated by blue-dashed box in Fig.~\ref{pspin}. This also implies that water may exist in all newly formed planets in our solar system. Magnetic field of planets may also play a dominate role in spin-spin coupling of large satellites at the time of their formation. The magnetic fields of the Sun and its planets may have evolved over the past 4.6 billion years, but the spin coupling among the planets remains strong and stable. On the other hand, orbital coupling or resonance in our solar system is governed by the gravitational influences between the parent star, the planets, and their satellites \citep[e.g.][]{gnp67, gre77}. \begin{figure} \epsscale{0.7} \plotone{f3.pdf} \caption[] {Schematic of expanding shocked gas regions produced by two reflected shock waves (SW$_i$$^\prime$ and SW$_2$$^\prime$) around a forming star. Rocky planets form in warm gas zone (WGZ), while gas giant planets form in cold gas zone (CGZ). \label{shockgas}} \end{figure} \begin{figure} \epsscale{0.9} \plotone{f4.pdf} \caption[] {Schematic of spin-spin coupling of planets in an external strong magnetic field of a younger star. The blue-dashed box shows the spin-spin coupling of eight planets in solar system. \label{pspin}} \end{figure} The initial conditions that determine the formation and final configuration of a star/planet system may vary from one group of dense clumps to another. Building such a stabilized system will strongly depend not only on the physical and chemical properties of its parent GMC and distribution of groups of dense clumps, but also on the triggering mechanisms for rapid fragmentation, core accretion, planet/satellite formation in the protoplanetary disk around a forming star. Therefore, each star/planet system may have its own unique characteristics depending on the initial conditions of the cloud, the clumps, and the ambient environment at the time of their formation. For example, recent detection of single planet systems found that HD 106906b (11 M$_{Jupiter}$) are located at a very large distance 650 AU to its parent star HD 106906 (K-type star) with only 13 million years age \citep{bai13}, and PSO J318.5-22 is a free-floating gas giant (6.5 M$_{Jupiter}$) formed without a parent star 12 million years ago \citep{mic13}. These two exoplanets may have formed directly from clump collapsing due to the onset of global gravitational instability in a cloud before they become fragmented. On the other hand, the HD 10180 is a multi-planet system; all of its planets are located at very close distances (0.02 - 3.5 AU) to their parent star HD 10180 (G1V star, 7.3 Gyr), but none of its planets has mass like gas giants (Jupiter or Saturn) \citep{tuo12}. The Gliese 876 system has four planets, two of them have Jupiter mass but are much closer to their parent star (0.1 - 0.2 AU) than those in our solar system \citep[and reference therein]{sha08}. The Gliese 876 system also has a notable orbital arrangement between its four planets, i.e. Laplace resonance configuration, which is similar to that of Jupiter's closest Galilean moons. On the other hand, four detected planets in the Upsilon Andromedae system are gas giant planets (0.6 - 4 M$_{Jupiter}$), located within 5 AU distance to their parent star \citep{hag08}. The star of the Upsilon Andromedae system is about 3.5 times more luminous than the Sun, and the distance of its inner most planet to its star is about six times closer than that of the Mercury to the Sun. Another multi-planet system that is similar to the Upsilon Andromedae system is the HR 8799 W system, but the inner most planet locates at a distance of 14 AU to its parent star that is 35 times further than that of the Mercury to the Sun \citep{opp13}. These extrasolar systems that have multiple planets may form in a similar way like our solar system, but have different clump properties, distribution, and protoplanetary disk conditions. From what has been said here, a hypothesis is an idea. More theoretical computation and observational studies are needed to draw more definite conclusions. These studies may include the properties of small groups of dense clumps and the interclump conditions in the GMCs, the role of shock waves, the interplay among thermal, dynamic and gravitational instabilities on large and small scales, their effects on planetary system formation and evolution, as well as observational measurements of gas and dust properties in nearby protoplanetary disks around forming/younger stars, e.g. the HD 142527 system \citep{fuk13}.
14
3
1403.0168
A new theoretical hypothesis on the origin and formation of the solar and extrasolar planetary systems is summarized and briefly discussed in the light of recent detections of extrasolar planets, and studies of shock wave interaction with molecular clouds, as well as H. Alfven's work on Sun's magnetic field and its effect on the formation of the solar system (1962). We propose that all objects in a planetary system originate from a small group of dense fragments in a giant molecular cloud (GMC). The mechanism of one or more shock waves, which propagate through the protoplanetary disk during the star formation is necessary to trigger rapid cascade fragmentation of dense clumps which in turn collapse quickly, simultaneously, and individually to form multi-planet and multi-satellite systems. Magnetic spin resonance may be the cause of the rotational directions of newly formed planets to couple and align in the strong magnetic field of a younger star.
false
[ "extrasolar planets", "molecular clouds", "dense fragments", "dense clumps", "rapid cascade fragmentation", "shock wave interaction", "-planet and multi-satellite systems", "recent detections", "formation", "Magnetic spin resonance", "the solar system", "Suns magnetic field", "a giant molecular cloud", "a planetary system", "the solar and extrasolar planetary systems", "the strong magnetic field", "GMC", "Sun", "turn", "multi" ]
9.115246
13.878332
-1
483101
[ "Teitler, Seth", "Königl, Arieh" ]
2014ApJ...786..139T
[ "Why is there a Dearth of Close-in Planets around Fast-rotating Stars?" ]
49
[ "National Academic Quiz Tournaments, LLC, 11521 W. 69th Street, Shawnee, KS 66203, USA", "Department of Astronomy &amp; Astrophysics and The Enrico Fermi Institute, The University of Chicago, Chicago, IL 60637, USA" ]
[ "2014ApJ...787..131Z", "2014ApJ...793L...3V", "2014ApJ...794..131L", "2014MNRAS.442.2081K", "2014MNRAS.443.1451L", "2014PASP..126..553P", "2014arXiv1410.3449A", "2015A&A...574A..39D", "2015ARA&A..53..409W", "2015ApJ...807...78F", "2015ApJ...809L..20M", "2015ApJ...813..101V", "2015MNRAS.446.3676A", "2015MNRAS.453.3720B", "2015csss...18..811L", "2016ApJ...816...18E", "2016ApJ...818....5L", "2016ApJ...820L...8M", "2016CeMDA.126..275B", "2016IAUFM..29A..57F", "2017A&A...604A.112G", "2017AJ....153...94C", "2017AJ....154....4P", "2017ApJ...844...12D", "2017ApJ...846L..13K", "2017ApJ...849L..11W", "2017MNRAS.465..149J", "2017MNRAS.466.1421Y", "2017MNRAS.471.2743T", "2018AJ....155..192T", "2018ApJ...864...65Q", "2018exha.book.....P", "2019A&A...621A.124B", "2019AJ....158..190H", "2019EAS....82....5M", "2019MNRAS.490.2390P", "2020AJ....159..150P", "2020ApJ...893L..29B", "2020IAUS..354..295A", "2021A&A...650A.126A", "2021ApJ...914...56A", "2022A&A...658A.195L", "2022A&A...666A.143K", "2022ASSL..466...85L", "2022ApJ...930L..23M", "2023A&A...679L..12G", "2023MNRAS.523.5220F", "2023NatAs...7..900C", "2024ApJ...960...50W" ]
[ "astronomy" ]
6
[ "planets and satellites: dynamical evolution and stability", "planet-star interactions", "stars: rotation", "Astrophysics - Earth and Planetary Astrophysics", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1981A&A....99..126H", "1997A&A...326.1023B", "1997ApJ...481..926M", "1998ApJ...500..428T", "2002A&A...386..222W", "2004ApJ...610..464D", "2007ApJ...661.1180O", "2009ApJ...698.1357J", "2009MNRAS.395.2268B", "2009MNRAS.396.1789P", "2010A&A...512A..77L", "2010ApJ...716.1269D", "2010ApJ...718L.145W", "2010ApJ...723.1583M", "2010ApJ...723L..64C", "2010ApJ...725.1995M", "2011ApJ...729...27B", "2011ApJ...733..115M", "2011ApJ...733L...9M", "2011ApJ...742...38Y", "2011ApJS..197....8L", "2011arXiv1109.2497M", "2012A&A...544A.124B", "2012ApJ...748L..10M", "2012ApJ...753..160W", "2012ApJ...753L..25R", "2012ApJS..201...15H", "2012MNRAS.425.2778M", "2013A&A...556A..36G", "2013ApJ...766...81F", "2013ApJ...775..105O", "2013ApJ...775L..11M", "2013ApJ...776....2L", "2013MNRAS.436.1883W", "2014A&A...565L...1P", "2014AAS...22315515W", "2014ApJS..210....1C" ]
[ "10.1088/0004-637X/786/2/139", "10.48550/arXiv.1403.5860" ]
1403
1403.5860_arXiv.txt
\label{sec:intro} The {\it Kepler}\/ space telescope has already led to tremendous advances in the detection and characterization of extrasolar planetary systems and in the study of the properties of their host stars. Until recently, these two aspects of the observed systems were considered separately, but the accumulation of data on both planets and stars is starting to make it possible to carry out statistical investigations of their joint characteristics. Such investigations can potentially shed new light on how these systems form and evolve, and, in particular, on how planets interact with their host stars. In a recent study, \citet[][hereafter MMA13]{McQuillanEtal13} derived the rotation period $P_{\rm rot}$ for 737 host stars of {\it Kepler}\/ Objects of Interest (KOIs) by applying an autocorrelation-function technique to the analysis of the observed star spot-induced photometric modulations. After plotting these periods against the orbital period $P_{\rm orb}$ of the innermost planet in each system, they noticed a clear dearth of close-in planets ($P_{\rm orb}\la 2-3\;$days) around rapidly rotating stars ($P_{\rm rot}\la 5-10\;$days). They fitted a line with a slope of $-0.69$ to the lower edge of the observed KOI distribution in the region of the $\log{P_{\rm orb}}-\log{P_{\rm rot}}$ plane bounded by $P_{\rm orb}\le 10\;$days and $P_{\rm rot} \ge 3 \;$days. They also pointed to the presence of several objects below the lower edge that exhibit near-synchronous rotation ($P_{\rm rot} \approx P_{\rm orb}$). These results were confirmed by \citet{WalkowiczBasri13}, who deduced the rotation periods of $\sim 950$ KOI hosts using the Fourier-based periodogram method. The latter authors found a compelling dearth of planets with $P_{\rm orb}$ and $P_{\rm rot}$ periods $< 6\;$days, and also noticed a concentration of planets with $P_{\rm rot} \approx P_{\rm orb}$. They furthermore pointed to the presence of a few systems with $P_{\rm rot} \approx 2\, P_{\rm orb}$, and noted that all the planets along these two loci have radii $R_{\rm p}>6\, R_\earth$. In this paper we propose that the short-period void in the $\log{P_{\rm orb}}-\log{P_{\rm rot}}$ plane of KOIs can be attributed to the tidal interaction between close-in planets and their (typically Sun-like) host stars, which, over the lifetimes of the observed systems, results in the spiraling-in of the nearest planets and the deposition of their orbital angular momenta in the host's envelope. Since the angular momentum of a planet of a given mass that moves on a Keplerian orbit scales as $P_{\rm orb}^{1/3}$ and its inspiral onset time is $\propto P_{\rm orb}^{13/3}$ (see Section~\ref{sec:formulate}), progressively lower values of $P_{\rm rot}$ correspond to the tidal ``ingestion'' of planets with progressively larger values of $P_{\rm orb}$. This naturally results in an inverse correlation between $P_{\rm rot}$ and the orbital period of the surviving closest-in planet. Using the inferred parameters of KOIs in a tidal interaction model, we demonstrate that this picture can indeed account for the lower edge of the empirical planet distribution. Our model also incorporates two physical processes that affect the stellar envelope and act to counter the spinup effect of planet ingestion. The first is core--envelope coupling, through which the envelope shares its angular momentum with the rest of the star. The second is magnetic braking, the process invoked to account for the apparent rotation--age correlation underlying the gyrochronology age-determination method for solar-type stars \citep[e.g.,][]{MeibomEtal11a,MeibomEtal11b}. The braking time is typically $<10^9\,$yr when $P_{\rm rot}<10\;$days, which explains our finding that the systems at the lower edge of the close-in planet distribution are generally among the youngest in the sample (see Section~\ref{sec:results}). The relative efficiency of this process would, however, inhibit the formation of synchronous systems, and we suggest (in line with previous work) that the observed $P_{\rm rot} \approx P_{\rm orb}$ locus likely corresponds to systems in which both magnetic braking and core--envelope coupling are weak. The possibility that the spin evolution of a star can be affected by the tidally induced ingestion of close-in planets was already considered previously in the literature. In particular, \citet{JacksonEtal09} suggested that this process could account for the observed orbital distribution of close-in planets and pointed out the dependence of the results on the initial $P_{\rm orb}$ distribution as well as on the systems' age distribution, \citet{Pont09} emphasized the dependence of the tidal interaction on the planet's mass, and \citet{BolmontEtal12} highlighted the potential implications of stellar spinups of this type to the reliability of the gyrochronology method. In this contribution we focus on the application of this idea to the distribution of low-period KOIs in the $\log{P_{\rm orb}}-\log{P_{\rm rot}}$ plane.
\label{sec:discuss} To determine the effect of the modeled tidal interaction on the adopted initial $P_{\rm orb}$ distribution, we obtained the probability density of the final distribution using Gaussian kernel density estimation. The resulting distribution for our fiducial case (Model~1) is shown by the solid blue curve in Figure~\ref{fig3}(b). It is noteworthy that, even though the initial distribution is a smooth power law in $P_{\rm orb}$, the final distribution manifests a break at $P_{\rm orb} \approx 4\;$days, where it steepens toward lower periods. The location of this break is consistent with the lower bound on the extent of the tidal interaction zone that we inferred from the properties of the final distribution in the $\log{P_{\rm orb}}-\log{P_{\rm rot}}$ plane for this case as well as with the boundary of the void that MMA13 identified in the observed distribution (see Section~\ref{sec:results}). We repeated the calculation for the cases shown in panels (a)--(c) of Figure~\ref{fig2} and determined that the result remains essentially the same when only the stellar spindown parameters are changed (Models~3 and~4), but that the break shifts to $P_{\rm orb} \approx 3\;$days when $Q^\prime_*$ is increased from $10^5$ to $10^6$ (Model~5; dashed red curve in Figure~\ref{fig3}(b)). Interestingly, the location of the latter break is approximately the same as in the data presented by \citet{Youdin11}, where it is fixed largely by the contribution of the $R_{\rm p} > 3\,R_\earth$ planets in his sample. We note in this connection that the values of the power-law index $\beta$ that characterize the asymptotic ``fast'' and ``slow'' regimes for each of the curves in Figure~\ref{fig3}(b) depend on our choice for the upper end of the $P_{\rm orb}$ interval ($P_{\rm orb,max}=10\;$days) and therefore cannot be directly compared with those given in \citet{Youdin11} (where $P_{\rm orb,max}=50\,$ days, and where the inferred values of $\beta$ also depend on the imposed slow/fast separation at $P_{\rm orb}=7\;$days). However, we verified that the presence of the break and its location do not depend on the value of $P_{\rm orb,max}$. These results indicate that the break is a robust feature of the final $P_{\rm orb}$ distribution that reflects the extent of the star--planet tidal interaction zone. The fact that the overall distribution is best represented by a model with $Q^\prime_* = 10^6$ even as its edge is better reproduced using $Q^\prime_* = 10^5$ is consistent with the possibility that this parameter evolves with stellar age (see Section~\ref{sec:results}). Our finding that the same basic model can account for the low-period structure of the KOI distribution in the $\log{P_{\rm orb}}-\log{P_{\rm rot}}$ plane and for the prominent break in the $P_{\rm orb}$ distribution of these objects lends strong support to the tidal interaction interpretation of {\it both}\/ the $P_{\rm orb}$ distribution of close-in planets {\it and} the $P_{\rm rot}$ distribution of their host stars.
14
3
1403.5860
We propose that the reported dearth of Kepler objects of interest (KOIs) with orbital periods P <SUB>orb</SUB> &lt;~ 2-3 days around stars with rotation periods P <SUB>rot</SUB> &lt;~ 5-10 days can be attributed to tidal ingestion of close-in planets by their host stars. We show that the planet distribution in this region of the log P <SUB>orb</SUB>-log P <SUB>rot</SUB> plane is qualitatively reproduced with a model that incorporates tidal interaction and magnetic braking as well as the dependence on the stellar core-envelope coupling timescale. We demonstrate the consistency of this scenario with the inferred break in the P <SUB>orb</SUB> distribution of close-in KOIs and point out a potentially testable prediction of this interpretation.
false
[ "rotation periods", "tidal interaction", "tidal ingestion", "orbital periods", "stars", "magnetic braking", "rot</SUB", "lt;~", "P <SUB", "KOIs", "the stellar core-envelope coupling timescale", "Kepler objects", "the log P <SUB>orb</SUB>-log P <SUB>rot</SUB> plane", "their host stars", "interest", "the P <SUB>orb</SUB> distribution", "the planet distribution", "Kepler", "close-in KOIs", ">rot</SUB" ]
8.060635
14.076192
-1
746070
[ "Momose, Rieko", "Ouchi, Masami", "Nakajima, Kimihiko", "Ono, Yoshiaki", "Shibuya, Takatoshi", "Shimasaku, Kazuhiro", "Yuma, Suraphong", "Mori, Masao", "Umemura, Masayuki" ]
2014MNRAS.442..110M
[ "Diffuse Lyα haloes around galaxies at z = 2.2-6.6: implications for galaxy formation and cosmic reionization" ]
135
[ "Institute for Cosmic Ray Research, The University of Tokyo, 5-1-5 Kashiwanoha, Kashiwa, Chiba 277-8583, Japan", "Institute for Cosmic Ray Research, The University of Tokyo, 5-1-5 Kashiwanoha, Kashiwa, Chiba 277-8583, Japan; Kavli Institute for the Physics and Mathematics of the Universe (WPI), The University of Tokyo, 5-1-5 Kashiwanoha, Kashiwa, Chiba 277-8583, Japan", "Kavli Institute for the Physics and Mathematics of the Universe (WPI), The University of Tokyo, 5-1-5 Kashiwanoha, Kashiwa, Chiba 277-8583, Japan; Department of Astronomy, Graduate School of Science, The University of Tokyo, 7-3-1 Hongo, Bunkyo-ku, Tokyo 113-0033, Japan", "Institute for Cosmic Ray Research, The University of Tokyo, 5-1-5 Kashiwanoha, Kashiwa, Chiba 277-8583, Japan", "Institute for Cosmic Ray Research, The University of Tokyo, 5-1-5 Kashiwanoha, Kashiwa, Chiba 277-8583, Japan; Center for Computational Sciences, The University of Tsukuba, 1-1-1 Tennodai, Tsukuba, Ibaraki 305-8577, Japan", "Department of Astronomy, Graduate School of Science, The University of Tokyo, 7-3-1 Hongo, Bunkyo-ku, Tokyo 113-0033, Japan; Research Center for the Early Universe, Graduate School of Science, The University of Tokyo, Tokyo 113-0033, Japan", "Institute for Cosmic Ray Research, The University of Tokyo, 5-1-5 Kashiwanoha, Kashiwa, Chiba 277-8583, Japan", "Center for Computational Sciences, The University of Tsukuba, 1-1-1 Tennodai, Tsukuba, Ibaraki 305-8577, Japan", "Center for Computational Sciences, The University of Tsukuba, 1-1-1 Tennodai, Tsukuba, Ibaraki 305-8577, Japan" ]
[ "2014ApJ...788...74S", "2014ApJ...789L..31W", "2014ApJ...792...34O", "2014ApJ...794..129H", "2014ApJ...795...33E", "2014PASA...31...40D", "2015A&A...573A..24C", "2015A&A...576A..51G", "2015ApJ...806...19D", "2015ApJ...806...46L", "2015ApJ...808...37C", "2015ApJ...809...19H", "2015ApJ...812..157H", "2015IAUS..309..295V", "2015MNRAS.451..400M", "2015MNRAS.453.3447D", "2016A&A...587A..98W", "2016A&ARv..24...14O", "2016ASSL..423..145D", "2016ApJ...816...16J", "2016ApJ...817...79H", "2016ApJ...818...38S", "2016ApJ...819...25K", "2016ApJ...822...84M", "2016ApJ...829....3A", "2016ApJ...830...52E", "2016ApJ...832..171T", "2016MNRAS.455.3991R", "2016MNRAS.456.4191P", "2016MNRAS.457.2318M", "2016MNRAS.457.3541C", "2016MNRAS.458..449M", "2016MNRAS.463.1678S", "2017A&A...606A..19G", "2017A&A...608A...6D", "2017A&A...608A...8L", "2017A&A...608A..10H", "2017ApJ...835..207G", "2017ApJ...837...11B", "2017ApJ...837..172X", "2017ApJ...838....4Y", "2017ApJ...839...17S", "2017ApJ...841...19M", "2017ApJ...843..133O", "2017ApJ...844..171Y", "2017ApJ...846...11M", "2017ApJ...846..134J", "2017IAUS..321..318K", "2017MNRAS.465..701J", "2017MNRAS.465.1543H", "2017MNRAS.466.1242S", "2017MNRAS.467.3306S", "2017MNRAS.469.2646C", "2017MNRAS.469.2959K", "2017MNRAS.471..267D", "2017MNRAS.472..772M", "2017arXiv170403416D", "2018A&A...616A..60O", "2018AJ....156...66M", "2018ApJ...862L..10E", "2018MNRAS.476.4725S", "2018MNRAS.476.5479P", "2018MNRAS.477.5406Y", "2018MNRAS.480.5140K", "2018MNRAS.481.1320C", "2018Natur.562..229W", "2018PASJ...70S..13O", "2019A&A...621A.107H", "2019A&A...623A.157S", "2019ApJ...871..164S", "2019ApJ...877....4L", "2019ApJ...878...68I", "2019ApJ...884....7E", "2019ApJ...887...85T", "2019ApJ...887..107F", "2019MNRAS.482.4744R", "2019MNRAS.484.2420C", "2019MNRAS.485.3947M", "2019MNRAS.488..120M", "2019MNRAS.489.2417A", "2019PASJ...71...55K", "2019SAAS...46....1D", "2019SAAS...46..189O", "2020A&A...635A..82L", "2020A&A...635A.154M", "2020ARA&A..58..617O", "2020ApJ...891..105W", "2020ApJ...891..177Z", "2020ApJ...901...41S", "2020MNRAS.492.1257M", "2020MNRAS.492.1778M", "2020MNRAS.493..141S", "2020MNRAS.498.3043M", "2020PhRvD.101h3032M", "2020arXiv200701322R", "2020arXiv201212378O", "2021A&A...650A..98W", "2021A&A...654A..80S", "2021ApJ...916...22K", "2021ApJ...921....4H", "2021MNRAS.506.5129B", "2021MNRAS.507.4294Z", "2021MNRAS.508...19C", "2022A&A...659A.183K", "2022A&A...660A..44K", "2022A&A...662A..64R", "2022A&A...666A..56M", "2022A&A...666A..78C", "2022A&A...668A.148R", "2022AJ....164..159P", "2022ApJ...926...31R", "2022ApJ...929...90L", "2022ApJ...931...97K", "2022ApJ...934...64A", "2022ApJS..262...38L", "2022MNRAS.516.1373F", "2023A&A...677A.125H", "2023ApJ...947...75K", "2023ApJ...953..118E", "2023ApJS..266...15M", "2023MNRAS.519..931W", "2023MNRAS.522.4275R", "2023MNRAS.523..676W", "2023MNRAS.523.5248B", "2023arXiv230905513G", "2023arXiv230906311G", "2024ApJ...961...63Z", "2024ApJ...962..175N", "2024ApJ...966..210L", "2024ApJS..271....5Z", "2024MNRAS.531.2701T", "2024PASA...41....5F", "2024arXiv240513122S", "2024arXiv240604399K", "2024arXiv240613915W" ]
[ "astronomy" ]
12
[ "galaxies: evolution", "galaxies: haloes", "galaxies: high-redshift", "dark ages", "reionization", "first stars", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
[ "1971PASP...83..199K", "1996PASP..108..699R", "2001ApJ...562..605F", "2002ApJ...571...40M", "2002PASJ...54..833M", "2004AJ....127..180C", "2004AJ....128.2073H", "2005ApJ...635L..13S", "2005MNRAS.363....2K", "2006AJ....132..117F", "2006ApJ...642L..13G", "2006ApJS..162....1G", "2006MNRAS.367..454H", "2006Natur.440..644M", "2007A&A...467..777C", "2007A&A...471..433P", "2007ApJ...657L..69L", "2007ApJ...666..663B", "2007ApJS..172....9T", "2007MNRAS.379.1599L", "2008ApJ...677...12O", "2008ApJ...681..856R", "2008ApJS..176....1F", "2008ApJS..176..301O", "2009AJ....138..923O", "2009ApJ...696..853L", "2009ApJ...703..785D", "2009ApJ...704.1640L", "2009ApJS..183..295T", "2009MNRAS.400..843G", "2009Natur.457..451D", "2010A&A...514A..64P", "2010ApJ...723..869O", "2010ApJ...724.1524O", "2010ApJ...725..633F", "2010MNRAS.402.1580O", "2010MNRAS.407..613G", "2011ApJ...729..140F", "2011ApJ...734..119K", "2011ApJ...736..160S", "2011ApJ...739...62Z", "2011MNRAS.415...11D", "2011MNRAS.415L...1G", "2012A&A...546A.111V", "2012ApJ...745...12N", "2012ApJ...750L..36M", "2012ApJ...752..114S", "2012MNRAS.423..344R", "2012MNRAS.424.1672D", "2012MNRAS.424.2193J", "2012MNRAS.425..878M", "2013ApJ...765...70H", "2013ApJ...765L..27H", "2013ApJ...769....3N", "2013ApJ...773..153J", "2013ApJ...776...75F", "2013ApJ...779...53Y", "2014ApJ...782....6H", "2014ApJ...785...64S", "2014ApJ...788...74S" ]
[ "10.1093/mnras/stu825", "10.48550/arXiv.1403.0732" ]
1403
1403.0732_arXiv.txt
\label{introduction} The circum-galactic medium (CGM) is closely related to galaxy formation and evolution. Gas inflows into galaxies could trigger starbursts (e.g., \citealp{dek09a,dek09b}), while gaseous outflows are thought to be a physical process of quenching star formation (e.g., \citealp{mori02,scan05,mori06,dave11}). The distribution of the CGM is characterized by Ly$\alpha$ emission, because Ly$\alpha$ photons escaping from a galaxy are resonantly scattered by surrounding neutral hydrogen gas. The scattered light would produce diffuse Ly$\alpha$ emission around a galaxy. This spatially-extended Ly$\alpha$ emission is dubbed Ly$\alpha$ halo (LAH). Numerical simulations have predicted that LAHs are ubiquitously present around high-$z$ galaxies (e.g., \citealp{lau07,zhen11,dijk12,verha12}). Radial surface brightness (SB) profiles of the LAHs are useful to understand kinematic properties of CGM and neutral hydrogen fraction of inter-galactic medium (IGM) at the epoch of cosmic reionization. \citet{zhen11} have predicted that a slope of radial SB profile depends on an outflowing velocity of CGM based on their radiative transfer model. \citet{dijk12} have calculated radiative transfer of Ly$\alpha$ photons propagating through clumpy and dusty large-scale outflowing inter-stellar medium (ISM), and reproduced an extended Ly$\alpha$ structure. Furthermore, \citet{jee12} demonstrate that radial SB profiles of LAHs are flatter at the epoch of reionization than at the post reionization epoch, due to Ly$\alpha$ photons scattered by neutral hydrogen of IGM. On the other hand, recent observations find no diffuse metal-line emission of hot ionized gas around high-$z$ galaxies on average \citep{yuma13}, and indicate that LAHs are probably not made by emission of hot CGM given by outflow, but the other physical processes of cold CGM. Extended Ly$\alpha$ emission has been observed around nearby star-forming galaxies (e.g., \citealp{ost09,hay13a,hay13b}) and QSOs (e.g., \citealp{rau08,goto09}). However, LAHs are too diffuse and faint to be detected for high-$z$ galaxies on an individual basis. LAHs at $z \geq 2$ are found in stacked data of $\sim20-2000$ narrowband (NB) images of high-$z$ galaxies in previous studies. \citet{haya04} have discovered a LAH around Lyman break galaxies (LBGs) at $z = 3.1$ with their composite NB image. \citet{stei11} have identified extended LAHs with a radius of $r \sim 80$ kpc around LBGs at a spectroscopic redshift of $\langle z\rangle=2.65$ by stacking $92$ NB images. \citet{matsu12} have stacked 130-864 LAEs at $z=3.1$, and detected LAHs. On the other hand, \citet{jian13} have found no extended Ly$\alpha$ emission in their composite image produced with dozens of LAEs at $z=5.7$ and $6.6$, although their results are based on the small statistics. There is an argument of systematic uncertainties producing spurious features similar to LAHs \citep{feld13}. \citet{feld13} have claimed that one of major sources of spurious LAHs is a large-scale point-spread function (PSF) that appears in deep images taken by ground-based observations \citep{king71}. A profile of large-scale PSF is largely extended, and the slope of profile changes at large radii of $>4\arcsec$ \citep{feld13}, probably due to atmospheric turbulence and instrumental conditions (e.g., \citealp{raci96,bern07}). The profile of large-scale PSF can mimic that of LAH, and would be mistakenly identified as an LAH. Thus, the existence of LAHs is still under debate. In order to test the existence of LAHs, a careful data analysis as well as a large galaxy sample is required. Here, we present our analysis and results of LAHs at $z=2.2-6.6$ based on our large LAE samples given by Subaru NB observations (\citealp{ouchi08,ouchi10,naka12}). The large LAE samples of high-quality Subaru images enable us to test the existence of diffuse LAHs and to extend the study from $z\sim 2-3$ to $6.6$. We show the data and analysis in Section \ref{data_and_analysis}, systematic errors in Section \ref{systematic_errors}, and our results of LAHs in Section \ref{results}, and discuss galaxy formation and reionization in Section \ref{discussions}. We summarize our results and discussions in Section \ref{summary}. Throughout this paper, we use AB magnitudes and adopt a cosmology parameter set of ($\Omega_m$, $\Omega_\Lambda$, $H_0$) = ($0.3$, $0.7$, $70$ km s$^{-1}$ Mpc$^{-1}$). In this cosmology, $1\arcsec$ corresponds to transverse sizes of ($8.3$, $7.6$, $7.2$, $5.9$, $5.4$) kpc at $z=$ ($2.2$, $3.1$, $3.7$, $5.7$, $6.6$). \begin{figure*} % \begin{center} \includegraphics[scale=0.25]{./fig1.eps} \end{center} \caption{ Composite continuum (top panels) and Ly$\alpha$ (bottom panels) images of our LAEs produced by the mean-combined method. From left to right panels, we show $z=2.2$, $3.1$, $3.7$, $5.7$, and $6.6$ LAE images. } \label{fig:image_mean} \end{figure*} \begin{figure*} \begin{center} % \includegraphics[scale=0.25]{./fig2.eps} \end{center} \caption{ Same as Figure \ref{fig:image_mean}, but for the median-combined method. } \label{fig:image_median} \end{figure*} \begin{figure*} \begin{center} % \includegraphics[scale=0.35]{./fig3.eps} \end{center} \caption{ Radial SB profiles of composite images of LAEs (solid lines) and PSFs (dotted lines) at redshifts of $z=2.2-6.6$. The upper and lower panels represent SB profiles of continuum and Ly$\alpha$ emission, respectively. The cyan and orange (blue and red) lines denote the results of mean-combined (median-combined) methods. } \label{fig:radi_allz} \end{figure*}
\label{discussions} \subsection{Do LAHs Really Exist?} The existence of LAHs around high-redshift galaxies is under debate \citep{feld13,jian13}. In Figure \ref{fig:radi_allz}, we identify statistically-significant extended Ly$\alpha$ emission around our LAEs at $z=2.2-6.6$ based on our unprecedentedly large LAE samples that allow us to achieve $\sim 10-100$ times deeper SB limits than those of typical previous studies (e.g., \citealt{stei11,feld13}; see Table \ref{tab:z}). In Section \ref{systematic_errors}, we examine potential systematic errors that would mimic LAHs, and find that the large-scale PSF of instrumental and atmospheric effects cannot produce radial profiles of our extended Ly$\alpha$ emission in SB and shape (Figure \ref{fig:lPSF}). Besides the large-scale PSF, there are a number of potential systematic effects, such as flat-fielding and sky subtraction errors (e.g., \citealt{feld13}) as well as unknown systematics. To reveal the total systematic errors involved in our data and analysis, we stack non-LAEs (Figure \ref{fig:sample_nonlae}) in the same manner as our LAEs, and carry out the empirical tests. We find that there exist systematic errors that make an extended emission signal, but that no systematic errors can make a radial profile with the SB amplitude and shape similar to those of our extended Ly$\alpha$ emission of LAEs, except for our sample of $z=3.7$ LAEs whose data quality is poor (Figure \ref{fig:radi_nonlae}). Thus, we conclude that we definitively identify LAHs around the LAEs with our data by our analysis technique. Our results indicate that LAEs commonly possess LAHs at $z=2.2-6.6$. \subsection{Physical Origin of LAHs} \label{physical_origin_of_lahs} In theoretical studies, LAHs are thought to be produced primarily by two physical mechanisms: 1) the resonant scattering of Ly$\alpha$ in the CGM and/or IGM (e.g., \citealp{lau07,zhen11,dijk12,verha12,jee12}), and 2) the cold streams (e.g., \citealp{fau10,goe10,ros12}). Our observational results imply that LAHs would be made by the former mechanism. First, we discuss the possibility of resonant scattering. Ly$\alpha$ photons escaping from a galaxy are scattered by the neutral hydrogen in the CGM within a few 100 kpc from the center of galaxy, making the Ly$\alpha$ emission extended more than that of stellar continuum. Theoretical predictions indicate that the SB profiles of LAHs are determined by a combination of the ISM dynamics and distribution (e.g., \citealp{lau07,zhen11,dijk12,verha12}). \citet{lau09a,lau09b} have carried out Monte Carlo radiative transfer simulations of Ly$\alpha$ propagating through the ISM with various kinematic properties, and found that the extent of LAHs is $r\sim50-100$ kpc. \citet{verha12} have predicted that the clumpy and inhomogeneous ISM produces an LAH with a characteristic radius of $r\sim10-20$ kpc in their cosmological simulations. The LAHs from our observations also extend up to a radius of $r\sim30-80$ kpc similar to those of the simulations (Figure \ref{fig:radi_allz}). Moreover, similar LAHs are predicted for the neutral IGM scattering Ly$\alpha$ photons at the epoch of reionization (e.g., \citealp{zhen11,jee12}). Thus, the sizes of observed LAHs are comparable to those of Ly$\alpha$ scattering models, which indicate that the major physical origin of the LAH is probably the resonant scattering of Ly$\alpha$ photons in neutral hydrogen of the CGM and/or IGM. Second, we examine whether the cold streams would be a major mechanism of the LAH formation. Cosmological hydrodynamic simulations have predicted that high-redshift ($z\geq2$) galaxies assemble baryon via accretion of relatively dense and cold gas ($\sim10^4$ K) that represents the cold streams (e.g., \citealp{keres05,dek09a,dek09b}). At the temperature of $\sim10^4$ K, the cold gas could primarily emit Ly$\alpha$ (e.g., \citealp{far01}), which would produce an LAH around galaxies \citep{fau10,goe10,ros12}. Cosmological simulations have predicted that Ly$\alpha$ emission are more extended for a higher dark halo mass of host galaxies. \citet{ros12} have found that cold streams in dark halos with a mass of $\geq10^{12}$ M$_\odot$ produce extended Ly$\alpha$ structures, but that Ly$\alpha$ emission is centrally-concentrated ($< 20$ kpc in diameter) in less-massive dark halos with a mass of $\sim10^{11}$ M$_\odot$. Because the typical dark halo mass of LAEs is estimated to be $\sim10^{11\pm1}$ M$_\odot$ (\citealp{ouchi10} and reference therein), the cold streams would not make an LAH as large as $30-80$ kpc found in our data. The cold streams are probably not the major mechanism of the LAH formation. \subsection{Scale-Length Dependence on Galaxy Population and Environment} \label{scale-length_dependence_on_galaxy_population_and_envrinment} In Section \ref{scale_lejgths_of_our_lahs_and_comparisons}, we find differences of scale lengths between measurements of this and some previous studies. First, at $z\sim 2$, the scale lengths from this study are smaller than those from \citet{stei11} by a factor of $\sim 2$ (Figure \ref{fig:size}). Note that the scale lengths of \citet{stei11} are estimated for their LBGs, while our study obtains the scale lengths of LAEs. By definition, Ly$\alpha$ emission of LAEs are brighter than that of LBGs, at a given SFR or UV continuum on average. Theoretical studies suggest that the resonance line of Ly$\alpha$ can escape from galaxies with a low column density of neutral hydrogen in the CGM \citep{verha12,zhen11}. This physical picture of Ly$\alpha$ escape is confirmed by recent imaging and spectroscopic observations \citep{shib14a,shib14b}. Based on this picture, Ly$\alpha$ photons produced in star-forming regions of LAEs reach the observer with little resonant scattering. The observer thus finds that Ly$\alpha$ profiles of LAEs are more centrally concentrated than those of LBGs, and obtains a bright Ly$\alpha$ luminosity in the central core of Ly$\alpha$ profile of LAEs above the shallow detection limit of non-stacked images. On the other hand, LBGs have more resonant scattering in the CGM than LAEs, and Ly$\alpha$ profiles of LBGs are largely extended. This is consistent with the fact that the SB of LBGs of \citealt{stei11} ($\sim 10^{-18}$ erg s$^{-1}$ cm$^{-2}$ Hz$^{-1}$ arcsec$^{-2}$ at 30 kpc) is one-order of magnitude brighter than that of our LAEs ($\sim 10^{-19}$ erg s$^{-1}$ cm$^{-2}$ Hz$^{-1}$ arcsec$^{-2}$ at 30 kpc). The scale-length difference of our LAEs and \citeauthor{stei11}'s LBGs indicates that LAEs do not have much neutral hydrogen CGM that scatters Ly$\alpha$, and that the Ly$\alpha$ profiles of LAEs are steeper than LBGs. Second, at $z \sim 3$, the LAH scale lengths of our LAEs are comparable with those of the lowest-density environment LAEs of \citet{matsu12}, but lower than the highest-density environment LAEs of \citet{matsu12} (Figure \ref{fig:size}). Here, the galaxy populations of our and \citeauthor{matsu12}'s samples are the same. The difference of scale lengths is probably explained by the environment. Clustering analysis with our LAE samples indicates that our survey field at $z=3.1$ is not a high density region \citep{ouchi08,ouchi10}. Our results confirm that low-density environment LAEs have a moderately small scale length of $r_n\simeq 5-10$ kpc, and support the idea of the environmental effect on LAHs that is claimed by \citet{matsu12}. Theoretical studies have predicted that the SB profiles are affected by the galaxy environment. \citet{zhen11} have found the SB profiles of their LAHs at $z=5.7$ have three notable features separated at two radial positions; a central cusp at $r\leq0.2$ Mpc (co-moving), a relatively-flat part at $r\sim0.2-1$ Mpc, and an outer steep region at $r\geq1$ Mpc. The two radial positions of $r=0.2$ and $1$ Mpc are characterized by the one- and two-halo terms of dark matter halos. However, we do not find such features in the radial profiles of our LAHs. This is probably because our LAHs at $z=5.7$ are only found within an inner region of $r\leq 0.2$ Mpc, and/or the clustering strength of the SXDS field is weaker than that of Zheng et al. (2011; private communication). One needs deeper images than those of this study to investigate the environmental effect at the moderately high redshift of $z=5.7$. \subsection{Size Evolution of LAHs} We find no size evolution of LAHs ($r_n\simeq 5-10$ kpc) from $z=2.2$ to $5.7$ as we describe in Section \ref{size_evolution_of_lahs}. \citet{mal12} have measured the half-light radii of stellar distribution of LAEs, $r_c$, in the rest-frame UV, and found no size evolution ($r_c\sim1$ kpc on average) in the redshift range of $z\simeq 2-6$. These two results indicate that the size ratio of LAHs to stellar component is almost constant, $r_n/r_c \sim 5-10$, between $z=2.2$ and $5.7$. This ratio is comparable with those of $z\sim 2$ LBGs ($r_n/r_c \sim5-10$; \citealt{stei11}) and $z\sim 0-3$ LAEs ($r_n/r_c \sim2-4$; \citealt{matsu12,hay13a}). The former LBG results indicate that there is a scaling relation between the LAH and stellar-distribution sizes over the samples of LBGs and LAEs. The latter should be comparable, because these results include a low-density environment LAE sample similar to ours. This no evolution of $r_n$ and $r_n/r_c$ at $z=2.2-5.7$ is interesting, because similar trends of no evolution of LAE's physical properties are found in this redshift range. Spectral energy distribution (SED) fitting of our LAE samples have revealed that stellar population of our LAEs do not evolve significantly in the redshift range of $z=2.2-5.7$ (\citealt{ono10a,ono10b,naka12}, see also \citealp{gaw06,pen07,pen10,fin11,mal12}). Moreover, hosting dark halos of LAEs are typically $M_{\rm DH}=10^{11\pm 1}$ M$_\odot$ at $z=2-7$ (\citealp{ouchi10} and reference therein), and do not evolve. These observational results of no evolution of $r_n$, $r_n/r_c$, stellar population, and dark halo mass would indicate that LAEs are the population in a specific stage of galaxy evolution, which NB observations can snapshot. We find a possible increase of scale length from $z=5.7$ to $6.6$ in Section \ref{size_evolution_of_lahs} (see Figure \ref{fig:size}). Again, clustering analysis with our LAEs indicates that our survey field at $z=6.6$ is not a high density region \citep{ouchi10}, and it is unlikely that this increase is due to the environmental effect. Because there is no significant increase of scale lengths in the redshift range of $z=2.2-5.7$, this sudden increase from $z=5.7$ to $6.6$ may be explained by cosmic reionization. Signatures of the increase in the neutral hydrogen fraction of IGM ($x_{\text{HI}}$) at $z\geq7$ have been found by many observational studies based on Ly$\alpha$ luminosity functions of LAEs or the Gunn-Perterson test of QSOs (e.g., \citealp{fan06,ota08,ouchi10,kashik11,goto11,shib12}). Numerical simulations show that Ly$\alpha$ emission around star-forming galaxies is extended due to the neutral hydrogen in the IGM (e.g., \citealp{zhen11,jee12}). \citet{jee12} have predicted that the SB profile of LAHs becomes flatter in the IGM with a high $x_{\text{HI}}$. The relatively large scale length of our LAHs at $z=6.6$ may come from a more neutral IGM at the epoch of reionization, although the reliability of increase is not very high in statistics (Figure \ref{fig:size}) and systematics (Section \ref{tests_for_all_systematic_errors}; see Section \ref{size_evolution_of_lahs} for the details). To conclude this trend, one needs a large amount of high-quality NB data such obtained by the upcoming survey of Subaru Hyper Suprime-Cam (HSC).
14
3
1403.0732
We present diffuse Lyα haloes (LAHs) identified in the composite Subaru narrow-band images of 100-3600 Lyα emitters (LAEs) at z = 2.2, 3.1, 3.7, 5.7, and 6.6. First, we carefully examine potential artefacts mimicking LAHs that include a large-scale point-spread function made by instrumental and atmospheric effects. Based on our critical test with composite images of non-LAE samples whose narrow-band-magnitude and source-size distributions are the same as our LAE samples, we confirm that no artefacts can produce a diffuse extended feature similar to our LAHs. After this test, we measure the scalelengths of exponential profile for the LAHs estimated from our z = 2.2-6.6 LAE samples of L<SUB>Lyα</SUB> ≳ 2 × 10<SUP>42</SUP> erg s<SUP>-1</SUP>. We obtain the scalelengths of ≃5-10 kpc at z = 2.2-5.7, and find no evolution of scalelengths in this redshift range beyond our measurement uncertainties. Combining this result and the previously known UV-continuum size evolution, we infer that the ratio of LAH to UV-continuum sizes is nearly constant at z = 2.2-5.7. The scalelength of our z = 6.6 LAH is larger than 5-10 kpc just beyond the error bar, which is a hint that the scalelengths of LAHs would increase from z = 5.7 to 6.6. If this increase is confirmed by future large surveys with significant improvements of statistical and systematical errors, this scalelength change at z ≳ 6 would be a signature of increasing fraction of neutral hydrogen scattering Lyα photons, due to cosmic reionization.
false
[ "non-LAE samples", "z", "neutral hydrogen scattering Lyα photons", "LAHs", "scalelengths", "cosmic reionization", "LAE", "diffuse Lyα haloes", "future large surveys", "increasing fraction", "composite images", "Lyα", "z ≳", "LAH", "significant improvements", "our LAE samples", "potential artefacts", "instrumental and atmospheric effects", "UV-continuum sizes", "exponential profile" ]
13.633943
6.450206
159
745677
[ "Traficante, A.", "Paladini, R.", "Compiegne, M.", "Alves, M. I. R.", "Cambrésy, L.", "Gibson, S. J.", "Tibbs, C. T.", "Noriega-Crespo, A.", "Molinari, S.", "Carey, S. J.", "Ingalls, J. G.", "Natoli, P.", "Davies, R. D.", "Davis, R. J.", "Dickinson, C.", "Fuller, G. A." ]
2014MNRAS.440.3588T
[ "The pros and cons of the inversion method approach to derive 3D dust emission properties in the ISM: the Hi-GAL field centred on (l, b) = (30°, 0°)" ]
3
[ "Jodrell Bank Centre for Astrophysics, School of Physics and Astronomy, The University of Manchester, Oxford Road, Manchester M13 9PL, UK; Dipartimento di Fisica, Università di Roma Tor Vergata, I-00133 Italy; Infrared Processing Analysis Center, California Institute of Technology, Pasadena, CA 91125, USA", "Infrared Processing Analysis Center, California Institute of Technology, Pasadena, CA 91125, USA", "Infrared Processing Analysis Center, California Institute of Technology, Pasadena, CA 91125, USA; Laboratoire d'Optique Atmospherique, UMR8518, CNRS-INSU, Université Lille 1, France", "Jodrell Bank Centre for Astrophysics, School of Physics and Astronomy, The University of Manchester, Oxford Road, Manchester M13 9PL, UK; Institut d'Astrophysique Spatiale, CNRS (UMR8617) Universite' Paris-Sud 11, F-91405 Orsay, France", "Observatoire astronomique de Strasbourg, Université de Strasbourg, CNRS, UMR 7550, France", "Department of Physics &amp; Astronomy, Western Kentucky University, Bowling Green, KY 42101, USA", "Infrared Processing Analysis Center, California Institute of Technology, Pasadena, CA 91125, USA", "Infrared Processing Analysis Center, California Institute of Technology, Pasadena, CA 91125, USA", "INAF, Istituto Fisica Spazio Interplanetario, I-00133 Rome, Italy", "Infrared Processing Analysis Center, California Institute of Technology, Pasadena, CA 91125, USA", "Infrared Processing Analysis Center, California Institute of Technology, Pasadena, CA 91125, USA", "Dipartimento di Fisica e Scienze della Terra, Università di Ferrara e Sezione INFN Ferrara, Via Saragat 1, I-44122 Ferrara, Italy; Agenzia Spaziale Italiana Science Data Center, c/o ESRIN, via Galileo Galilei, I-00044 Frascati, Italy; INAF/IASF Bologna, Via Gobetti 101, I-40129 Bologna, Italy", "Jodrell Bank Centre for Astrophysics, School of Physics and Astronomy, The University of Manchester, Oxford Road, Manchester M13 9PL, UK", "Jodrell Bank Centre for Astrophysics, School of Physics and Astronomy, The University of Manchester, Oxford Road, Manchester M13 9PL, UK", "Jodrell Bank Centre for Astrophysics, School of Physics and Astronomy, The University of Manchester, Oxford Road, Manchester M13 9PL, UK", "Jodrell Bank Centre for Astrophysics, School of Physics and Astronomy, The University of Manchester, Oxford Road, Manchester M13 9PL, UK" ]
[ "2015MNRAS.450.2025A", "2015MNRAS.450.4277V", "2016A&A...591A.149M" ]
[ "astronomy" ]
7
[ "dust", "extinction", "Galaxy: structure", "infrared: ISM", "Astrophysics - Galaxy Astrophysics" ]
[ "1967ApJ...147..471M", "1970A&A.....6..364W", "1974A&AS...17..251W", "1978A&A....66...65S", "1978AuJPA..45....1H", "1983A&A...128..212M", "1983MNRAS.204...53S", "1985ApJ...288..618R", "1986A&A...154...25B", "1986ApJS...60....1S", "1988ApJ...330..964B", "1989ApJ...342..272F", "1990A&A...233..437B", "1990ARA&A..28..215D", "1991ARA&A..29..195C", "1994A&A...286..203G", "1996A&A...312..256B", "1996PASA...13..243S", "1997ApJ...480..173S", "2000A&A...354..247L", "2000A&AS..147..335D", "2001RvMP...73.1031F", "2003A&A...397..133R", "2003A&A...409..523R", "2003ApJ...586.1067H", "2003MNRAS.341..369D", "2003PASP..115..953B", "2004A&A...422L..47S", "2004ASSL..315...47G", "2004ApJ...603..560S", "2004MNRAS.347..237P", "2005ApJ...626..195G", "2005ApJ...631.1001K", "2005MNRAS.360.1545P", "2005PASP..117.1113M", "2005Sci...307.1292G", "2006A&A...453..635M", "2006AJ....131.1163S", "2006AJ....132.1158S", "2006ApJ...653.1226Q", "2006ApJS..163..145J", "2007A&A...465..839P", "2007ApJ...660..346P", "2007MNRAS.379.1599L", "2008AJ....136..919B", "2008ApJ...679..481P", "2009A&A...505..405P", "2009ApJ...690..706A", "2009ApJ...707..510Z", "2009ApJS..182..131R", "2009MNRAS.400.2050G", "2009PASP..121...76C", "2009PASP..121..213C", "2009RvMP...81..969H", "2010A&A...518L...2P", "2010A&A...518L...3G", "2010A&A...518L..88B", "2010A&A...518L..90B", "2010A&A...518L.101Z", "2010A&A...520L...8P", "2010ASPC..438..111G", "2010ApJ...710..133A", "2010ApJ...716.1191W", "2010ApJ...721..686P", "2010ApJ...723..492R", "2010ApJ...724L..44C", "2010MNRAS.405.1654A", "2010PASP..122..314M", "2011A&A...525A.103C", "2011A&A...527A.141C", "2011A&A...535A.128B", "2011A&A...536A..19P", "2011A&A...536A..20P", "2011A&A...536A..21P", "2011A&A...536A..22P", "2011A&A...536A..23P", "2011A&A...536A..24P", "2011A&A...536A..25P", "2011ApJ...732..100D", "2011ApJ...735....6P", "2011ApJ...741..110D", "2011MNRAS.416.2932T", "2012ApJ...754...94T", "2012ApJ...760..149P", "2012MNRAS.422.2429A", "2013A&A...554A.103P", "2013ARA&A..51..207B" ]
[ "10.1093/mnras/stu493", "10.48550/arXiv.1403.3327" ]
1403
1403.3327_arXiv.txt
A variety of space-borne experiments in the course of the last two decades (IRAS, Spitzer, \textit{Herschel}) have shown that the Galaxy is filled with relatively cold dust (15$\leq\mathrm{T}_{d}\leq30$ K, e.g. \citet{Ferriere01} distributed along each Line of Sight (LOS), and associated with the atomic, molecular and ionized gas phases. This finding has also been recently confirmed by the \textit{Planck} whole-sky observations \citep{Planck_A22,Planck_A23,Planck_A24,Planck_A25}. Dust grains can absorb and re-emit a large fraction of the radiation provided by stellar sources. Depending on the size, they are stochastically heated (i.e. the smaller grains, which emit mainly in the Near/Mid-IR), or reach thermal equilibrium (i.e. the bigger grains, which emit predominantly in the Far-IR regime). Dust emission properties have been intensely studied at high Galactic latitudes, where each gas phase can be assumed relatively isolated from the others, and mixing along the LOS can be avoided \citep[e.g.][]{Boulanger88, Boulanger96,Lagache00}. On the contrary, in the Galactic Plane disentangling dust emission arising from different gas components and intrinsically different environments is a complicated problem which requires additional kinematic information on the emitting gas \citep[e.g.][]{Bloemen90}. The separation can be achieved with the so-called \textit{inversion method}, originally introduced by \citet{Bloemen86} to analyse the individual contribution of atomic and molecular gas to the integrated $\gamma$-ray emission. The model was later extended by \citet{Bloemen90} to include the FIR emission across the Galactic Plane observed with IRAS at 60 \um\ and 100 \um . So far, the application of the inversion method has been limited by the angular resolution of the available IR and ancillary data (equal or close to $1^{\circ}$) which does not allow to, e.g., spatially separate individual clouds \citep{Bloemen90, Giard94, Sodroski97, Paladini07, Planck_Marshall11}. These works have demonstrated that, on average, most of the Galactic IR luminosity is associated with dust in the atomic gas component, which is primarily irradiated by the local radiation field. Dust associated with the molecular and ionized components, on the other hand, is mostly heated by O and B stars embedded in molecular clouds \citep{Sodroski97,Paladini07}. \citet{Planck_Marshall11} decomposed the emission from 12 $\mu$m to millimeter wavelengths using IRAS, COBE-DIRBE, WMAP and the recent \textit{Planck} data in the latitude range $\vert b\vert\leq10^{\circ}$ and found evidence of the existence of a significant amount of cold atomic and warm molecular gas not accounted for by the standard tracers. This gas is typically referred to as \textit{dark gas} \citep{Grenier05}. In particular, the authors of this work claim that the \textit{dark gas} column density is comparable - or greater - to the column density of molecular gas outside the Molecular Ring, i.e. a region of the Galaxy comprised between Galactocentric radius 4 kpc $<$ R $<$ 8 kpc, where 70 percent of the molecular gas resides \citep{Combes91}. As well as using low-spatial resolution, early inversion works often did not include the ionized gas component, as historically its contribution to the overall IR emission was thought to be negligible \citep{Bloemen90,Giard94}. A few studies \citep{Sodroski97, Paladini07, Planck_Marshall11} did take this phase of the gas into account, but used the only available data at the time for tracing ionized gas, that is radio continuum data. There are two problems with this approach. The first is that radio continuum emission is not uniquely associated with the interaction - and deceleration - of free electrons with ions in a plasma, the free-free or bremsstrahlung emission: at low frequencies (5 GHz or less), one has to estimate and subtract a possible contamination due to synchrotron emission \citep[e.g.][]{Paladini05}, while at relatively high frequencies ($>$ 10 GHz), spinning dust emission may become significant \citep[e.g.][]{Planck_Dickinson11, Planck_Marshall11}. The second, even more important, issue is the fact that radio continuum data are unable to provide the 3D-information on the location of the emitting gas which is required by inversion techniques. The work we describe in this paper is therefore motivated by the following considerations: \begin{enumerate} \item the resolution and sensitivity of the newly available Spitzer and \textit{Herschel} data dramatically improve our view of the Galactic Plane at IR wavelengths: the combined GLIMPSE \citep{Benjamin03}, MIPSGAL \citep{Carey09} and Hi-GAL \citep{Molinari10_PASP} surveys have mapped the inner Galactic Plane in the wavelength range 8 $\mu$m $\leq\lambda\leq$ 500 $\mu$m with a resolution of 35 arcsec, or higher. These new data allow us to set more stringent constraints on dust properties and on their variations with Galactocentric radius; \item the last couple of years have witnessed a tremendous improvement in the quality of available data on the ionized gas. In particular, hydrogen recombination lines (RRLs) have been observed for large portions of the Galactic Plane \citep{Staveley-Smith96,Alves11}. These data are sampled with a resolution of $\simeq15$ arcmin, which allows us to work with a $\sim$4 times better resolution than the previous inversions. Even more important, they provide unprecedented information about the properties and distribution of the ionized gas component along the LOS; \item in previous inversion works, the decomposition into Galactocentric bins has been done {\em{blindly}}, that is without taking into account local features present at a given location of the Galaxy or on specific angular scales. The higher resolution of the IR as well as of the ancillary data allow now a more targeted approach; \item last but not least, if on one side the higher angular resolution of the available data makes it possible to devote more attention to the specific content of the Galactic region to {\em{invert}}, on the other hand it sheds light on the limitations of the inversion technique itself. Hence, one of the goals of this work is to investigate and describe these limitations. \end{enumerate} In this first paper we concentrate on a $2^{\circ}\times2^{\circ}$ field centred on \textit{(l,b)}=(30$^{\circ}$.0,0$^{\circ}$.0) observed by the GLIMPSE, MIPSGAL and Hi-GAL surveys. This field was one of two observed during the Hi-GAL Science Demonstration Phase (SDP). Therefore, hereafter we will refer to it as SDPF1 ({\em{Science Demonstration Phase Field 1}}). The paper is organized as follows. A review of the content of SDPF1 is presented in Section \ref{sec:l30_field}. The IR and ancillary radio data used for the analysis are described in Section \ref{sec:infrared_ancillary}. The inversion model is discussed in Section \ref{sec:model_data}, as well as the Galactocentric region subdivision and the gas column density evaluation for each gas phase. We find evidence of untraced cold atomic and possibly warm molecular gas features extending up to several arcmins. These features do not allow a correct evaluation of the gas column densities, thereby preventing the inversion model from working properly. The regions where these features are dominant, however, can be predicted by analyzing the correlation between the gas column densities and the IR maps, as explained in Section \ref{sec:pearsons}. In Section \ref{sec:results_discussion} we present our results, and discuss the limitations of the inversion model in its current stage of development. We also investigate the importance of accounting for dust associated with the ionized gas, by demonstrating with a simple test the erroneous conclusions that one may reach by neglecting this component. In Section \ref{sec:conclusions} we describe our conclusions.
\label{sec:conclusions} We have investigated dust properties in a 2$\times$2 square degrees Hi-GAL field (SDPF1) centred on (\textit{l,b})=(30$^{\circ}$,0$^{\circ}$), in the wavelength range 8 \um\ $\leq\lambda\leq500$ \um. For this purpose, we have used an inversion technique, first introduced by \citet{Bloemen86}, to decompose the observed integrated IR emission into individual contributions associated with dust in the atomic, molecular and ionized phase of the gas and located at different Galactocentric distances. We have used, for the first time in an inversion analysis, Radio Recombination Lines (RRLs) to trace the ionized gas. In addition, the decomposition into Galactocentric bins (or {\em{rings}}) is performed exploiting the natural boundaries of the structures (i.e. segments of spiral arms) as they appear in the gas data cubes. We have solved the inversion equation for all the decomposition rings (i.e. Ring 1, 2, 3 and 4), and obtained positive solutions only for Ring 1. A Pearson's coefficient and longitude profiles analyses reveal that Ring 1, which covers Galactocentric distances between 4.2 and 5.6 kpc and hosts the mini-starburst W43, dominates the total IR emission towards SDPF1. For this ring only, we have fitted with DustEM the emissivities retrieved by the inversion method. These fits allow estimating, for each phase of the gas, dust temperatures and abundances, as well as the intensity of the local radiation field normalized to the intensity of the radiation field in the solar neighbourhood. In particular we find, for the ionized gas phase with respect to the other gas phases, an indication of PAH depletion and an intensity of the local radiation field two times higher, which reflects into a higher dust temperature. For the other rings (Ring 2, 3 and 4), the inversion equation either cannot be solved or returns negative emissivities. The Pearson's coefficients suggest a weak degree of correlation with the IR templates and, in a few cases, even an anti-correlation. For Ring 2 and 3, this result might be ascribed to the presence of a large amount of \textit{untraced gas}, either associated with warm \H2 and/or cold HI. This hypothesis could find support in the fact that, in Ring 3, cold HI structures are indeed found. In this scenario, the column densities derived from the standard tracers would not be able to fully account for the observed IR emission, hence the assumption of the inversion model would break down and the resulting emissivities (e.g. negative) be unreliable. In Ring 4, which covers the outer Galaxy, the slight degree of anti-correlation with the input IR maps is probably indicative of the intrinsic low emissivity of the region, due to a combined drastic decrease of both hydrogen column density and intensity of the interstellar radiation field. We have investigated the role of extinction in evaluating total column densities along the LOS as an alternative method to gas tracers. For this purpose, we have attempted to build an extinction map for SDPF1 in two independent ways, i.e. using deep catalogues of giant stars and with a 3D-extinction model. Although both methods appear to be promising, they currently face the severe limitation of not being able to trace extinction beyond $\sim$ 10 -- 15 kpc from the Sun. Finally, we have explored the impact of neglecting the ionized gas phase in inversion analysis, as often done in the past. We have shown that, by not including this gas component, the temperature of dust associated with the molecular phase is artificially increased, due to its high degree of correlation with the input IR templates. We conclude with a general remark. In this work we have improved, with respect to previous inversion studies, the estimation of the HI, \H2 and HII column densities. However, we believe that this analysis has shown, above all, that local effects, such as departures from circular motion and the presence of cold HI structures (and, likely, warm \H2), become important when a 3D-inversion is performed in small sections of the Plane and with an angular resolution higher or comparable to the angular scale on which these effects dominate the total emission. Conversely, when the entire Galactic Plane is {\em{inverted}} at low angular resolution, the peculiarities of each LOS are averaged out. Further developments of 3D-inversion models will have to take into account these limitations, both by including non-radial motions, i.e. the grand spiral design of the Galaxy, and by estimating total column densities accounting for the blending of cold and warm material along the LOS.
14
3
1403.3327
Herschel far-infrared continuum data obtained as part of the Hi-GAL survey have been used, together with the GLIMPSE 8 μm and MIPSGAL 24 μm data, to attempt the first 3D-decomposition of dust emission associated with atomic, molecular and ionized gas at 15 arcmin angular resolution. Our initial test case is a 2 × 2 square degrees region centred on (l, b) = (30°, 0°), a direction that encompasses the origin point of the Scutum-Crux Arm at the tip of the Galactic Bar. Coupling the IR maps with velocity maps specific for different gas phases (H I 21cm, <SUP>12</SUP>CO and <SUP>13</SUP>CO, and radio recombination lines), we estimate the properties of dust blended with each of the gas components and at different Galactocentric distances along the line of sight (LOS). A statistical Pearson's coefficients analysis is used to study the correlation between the column densities estimated for each gas component and the intensity of the IR emission. This analysis provides evidence that the 2 × 2 square degree field under consideration is characterized by the presence of a gas component not accounted for by the standard tracers, possibly associated with warm H<SUB>2</SUB> and cold H I. We demonstrate that the IR radiation in the range 8 &lt; λ &lt; 500 μm is systematically dominated by emission originating within the Scutum-Crux Arm. By applying an inversion method, we recover the dust emissivities associated with atomic, molecular and ionized gas. Using the DustEM model, we fit the spectral energy distributions for each gas phase, and find average dust temperatures of T<SUB>d,H I</SUB> = 18.82 ± 0.47 K, T<SUB>d,H<SUB>2</SUB></SUB> = 18.84 ± 1.06 K and T<SUB>d,H II</SUB> = 22.56 ± 0.64 K, respectively. We also obtain an indication for polycyclic aromatic hydrocarbons depletion in the diffuse ionized gas. We demonstrate the importance of including the ionized component in 3D-decompositions of the total IR emission. However, the main goal of this work is to discuss the impact of the missing column density associated with the dark gas component on the accurate evaluation of the dust properties, and to shed light on the limitations of the inversion method approach when this is applied to a small section of the Galactic plane and when the working resolution allows sufficient de-blending of the gas components along the LOS.
false
[ "different gas phases", "ionized gas", "dust emission", "average dust temperatures", "dust", "the dark gas component", "emission", "a gas component", "each gas component", "the gas components", "LOS", "IR", "different Galactocentric distances", "Galactic", "<SUP>13</SUP", "each gas phase", "atomic, molecular and ionized gas", "<", "CO", "the total IR emission" ]
11.96417
9.270433
-1
2053970
[ "Watson, Douglas F.", "Hearin, Andrew P.", "Berlind, Andreas A.", "Becker, Matthew R.", "Behroozi, Peter S.", "Skibba, Ramin A.", "Reyes, Reinabelle", "Zentner, Andrew R.", "van den Bosch, Frank C." ]
2015MNRAS.446..651W
[ "Predicting galaxy star formation rates via the co-evolution of galaxies and haloes" ]
56
[ "Kavli Institute for Cosmological Physics, 5640 South Ellis Avenue, The University of Chicago, Chicago, IL, USA;", "Kavli Institute for Cosmological Physics, 5640 South Ellis Avenue, The University of Chicago, Chicago, IL, USA; Fermilab Center for Particle Astrophysics, Fermi National Accelerator Laboratory, Batavia, IL, USA; Yale Center for Astronomy &amp; Astrophysics, Yale University, New Haven, CT, USA", "Department of Physics and Astronomy, Vanderbilt University, Nashville, TN, USA", "SLAC National Accelerator Laboratory, Menlo Park, CA 94025, USA; Kavli Institute for Particle Astrophysics and Cosmology, Stanford, CA 94309, USA; Department of Physics, School of Humanities and Sciences, Stanford University, Stanford, CA 94309, USA", "Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21218, USA", "Department of Physics, Center for Astrophysics and Space Sciences, University of California, 9500 Gilman Dr., La Jolla, San Diego, CA 92093, USA", "Kavli Institute for Cosmological Physics, 5640 South Ellis Avenue, The University of Chicago, Chicago, IL, USA", "Department of Physics and Astronomy, University of Pittsburgh, Pittsburgh, PA 15260, USA; Pittsburgh Particle physics, Astrophysics and Cosmology Center (PITT PACC), USA", "Department of Astronomy, Yale University, PO Box 208101, New Haven, CT, USA" ]
[ "2014ApJ...787..156B", "2014MNRAS.444..729H", "2014MNRAS.445.1713V", "2015ApJ...803...26P", "2015ApJ...807..152S", "2015ApJ...808...40W", "2015MNRAS.446.1939F", "2015MNRAS.450.1546B", "2015MNRAS.451L..45H", "2015MNRAS.452..444C", "2015MNRAS.452.1958H", "2015MNRAS.452.4013K", "2015MNRAS.454.3020B", "2015MNRAS.454.3030P", "2015arXiv150703605B", "2016ApJ...819..119L", "2016MNRAS.455..158V", "2016MNRAS.455..185B", "2016MNRAS.456.1459M", "2016MNRAS.456.4364B", "2016MNRAS.459.1929T", "2016MNRAS.460.1457S", "2016MNRAS.460.2552H", "2016MNRAS.461.2135H", "2016MNRAS.461.3421F", "2016MNRAS.463.1907F", "2017ApJ...838...87C", "2018ARA&A..56..435W", "2018ApJ...853...84Z", "2018ApJ...857..127S", "2018MNRAS.477..359C", "2018MNRAS.478.4487T", "2019ApJ...884...76B", "2019ApJ...887...17Z", "2019ApJ...887...40W", "2019MNRAS.488.3143B", "2019MNRAS.489.1805H", "2019OJAp....2E...7A", "2021AJ....161...49B", "2021ApJ...906...97F", "2021MNRAS.506.3155J", "2021MNRAS.507.5320Z", "2021OJAp....4E...7H", "2022ApJ...931..145W", "2022ApJ...940...13D", "2022MNRAS.511.4900S", "2022OJAp....5E...3H", "2023MNRAS.518..562A", "2023MNRAS.519.1648A", "2023MNRAS.519.4253L", "2023MNRAS.520.1774W", "2023MNRAS.522.3188W", "2023MNRAS.523.6386M", "2023arXiv230209199L", "2024JCAP...02..048F", "2024arXiv240201834H" ]
[ "astronomy" ]
7
[ "galaxies: evolution", "galaxies: haloes", "galaxies: star formation", "cosmology: theory", "dark matter", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
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[ "10.1093/mnras/stu2065", "10.48550/arXiv.1403.1578" ]
1403
1403.1578_arXiv.txt
\label{sec:intro} One of the principal goals of galaxy evolution theory is to understand the connection between the properties of galaxies and their host dark matter halos. There is now a well-established relation between the stellar mass of a galaxy and the mass of the halo in which it resides \citep[e.g.,][]{yang12,leitner12,wang_etal12,moster13,behroozi13b,kravtsov_size_Rvir13,kravtsov_smhm_14}. Moreover, the fact that the stellar mass-to-halo mass connection remains tight across cosmic time \citep{conroy_wechsler09,behroozi13,watson_conroy13,lu_etal14} suggests that there are likely further links between halo properties and the star formation rate (SFR) of galaxies. With this as motivation, the aim of the present work is to address the following question: is there a simple link between the SFR of galaxies and the dark side of the universe? The complex nature of star formation in galaxies indicates that the relationship between the SFR of a galaxy and the properties of its host dark matter halo may be complicated. First, at {\em fixed luminosity} or {\em fixed stellar mass}, there exists a clear bimodality in the distribution of galaxy color/SFR, with distinct red/quenched and blue/star-forming populations \citep{blanton03,baldry04,bell04,blanton05,cooper06,wyder07,wetzel_etal11,cooper12}. Additionally, galaxy color/SFR depends on environment: denser environments, such as rich groups and clusters, are populated by significantly more red sequence galaxies than actively star-forming ones \citep{balogh99,blanton05,weinmann06b,weinmann09,peng_etal10,peng_etal12,carollo_etal12,tal_etal14}. Furthermore, the specific processes that attenuate SFR in a 'central' galaxy (the galaxy at the minimum of the halo potential well) may be distinct from those governing the 'satellite' galaxies orbiting the central \citep{vdbosch_08}. Finally, the SFR/color dependence of galaxy location within the cosmic web also manifests in measurements of two-point statistics; as a function luminosity or stellar mass, red/quenched galaxies exhibit stronger clustering than blue/star forming galaxies \citep[e.g.,][]{norberg02,zehavi02,zehavi05a,li06,zehavi11,yang12,mostek12,guo_etal14}. Such observed complexities may lead one to conclude that complicated modeling of the physics governing the quenching of galaxies is required to reproduce observed galaxy statistics. However, in a pair of recent papers introducing the {\em age matching} formalism \citep[][hereafter Papers I $\&$ II, respectively]{HW13a,HW13b}, it was shown that in fact a very simple model for galaxy color can account for the rich variety of trends exhibited by galaxies in the low-redshift universe. The central hypothesis of age matching is that at fixed luminosity (or fixed stellar mass), galaxy color is in monotonic correspondence with a proxy for halo age, at fixed halo maximum circular velocity $\vmax$. In Paper I, this prescription was shown to accurately reproduce the observed $g-r$ color-dependent clustering of galaxies in the Sloan Digital Sky Survey \citep[SDSS:][]{york00a} for the luminosity-selected galaxy samples presented in \citet{zehavi11}, as well as the scaling between $g-r$ color and host halo mass. In Paper II, similar success of the age matching formalism was demonstrated for model predictions of new measurements of both SDSS clustering and galaxy-galaxy lensing as a function of stellar mass and $g-r$ color. However, $g-r$ color and star formation activity are not perfectly correlated. For instance, galaxies that are actively forming stars can often appear red due to the ubiquitous presence of dust \citep[e.g.,][]{stein_soifer83,maller_etal09,masters_etal10}. Furthermore, $g-r$ color is the convolution of many physical properties of galaxies, including: stellar age, metallicity, and instantaneous SFR. Thus, a model for the color-dependence of galaxy location within the cosmic web may not smoothly translate to a model for the SFR dependence. In the present study, we demonstrate how age matching, without modification to the technique introduced in Papers I $\&$ II, is equally successful at reproducing new SDSS measurements of stellar mass- and SFR-dependent clustering and galaxy-galaxy lensing. As we will demonstrate in a forthcoming paper (Watson, Skibba $\&$ Hearin 2014, in prep) studying {\em marked} correlation functions \citep{skibba06,skibbamarkedCF13}, this simultaneous success of our model is primarily due to the surprising observational fact that the two-point function is almost entirely insensitive to the choice of SFR or $g-r$ as a star formation indicator. Additionally, in this paper we take a sharp focus on the population of satellite galaxies. While satellites are in the minority by number, the physics governing satellite galaxy SFR is a key ingredient to painting a complete picture of the theory of galaxy evolution. Satellite galaxies can be subject to a number of complex processes which are believed to stifle their star formation as they orbit within the gravitational potential well of their host halo. These include the removal of cold gas from the disc due to ram pressure \citep{gunn_gott72}, the stripping of the surrounding hot gas reservoir, known as `strangulation' \citep{larson80}, disruption of satellite galaxies due to tidal stripping \citep{purcell_etal07, watson_etal12b}, and `harassment' by gravitational interactions with other nearby galaxies \citep{moore_etal98}. In age matching there is {\em no} explicit modeling of such post-accretion processes. And yet, we will demonstrate that this remarkably simple model accurately predicts the radial profiles of star-forming and quenched satellite galaxies within and around the environment of groups, rich groups, and clusters. As discussed in \S~\ref{sec:discussion}, the success of age matching at reproducing these trends indicates that in much of the literature on satellite evolution, the influence of post-accretion processes on quenching satellites has been over-estimated. The paper is laid out as follows. In \S~\ref{sec:data} we describe the data, simulation and halo catalogs incorporated throughout this work. An overview of the age matching and the more generic ``conditional abundance matching'' formalism is given in \S~\ref{sec:model}. In \S~\ref{sec:results} we present our main results. Specifically, in \S~\ref{subsec:2PCF_gg_results} we show our model predictions for new measurements of the SFR-dependent two-point correlation function and galaxy-galaxy lensing signal. In \S~\ref{subsec:sat_results} we study the spatial properties of star-forming and quenched satellite galaxies within and around halos. In \S~\ref{sec:discussion} we provide a discussion and interpretation of our findings. We conclude in \S~\ref{sec:summary} with a brief summary of our primary results. Throughout this work we assume a flat $\Lambda$CDM cosmological model with $\Omega_{\mathrm{m}}=0.27$ and Hubble constant $H_0=70$ km s$^{-1}$ Mpc$^{-1}$. \begin{figure*} \begin{center} \includegraphics[width=1.\textwidth]{./FIGS/ssfr.pure.PDFs.eps} \caption{ The probability distribution functions (PDFs) of the specific star formation rate (sSFR) of galaxies in our mock catalog (gray solid curves) as compared to those measured in the SDSS galaxy catalogs (dotted black curves). By construction, the PDFs of sSFR of our mock galaxies are in exact agreement with the data for all the galaxies in our sample (top left panel) as well as three stellar mass threshold samples (bottom row): $\log_{10}(\Mstar)>[9.8,10.2,10.6]$. Since our sSFR assignment to mock galaxies is blind to the distinction between central and satellite galaxies, the resultant PDFs in the center and right panels of the top row are {\em predictions} of age matching, and demonstrate a non-trivial success of the technique.} \label{fig:sm_SFR_PDFs} \end{center} \end{figure*}
\label{sec:discussion} \subsection{Simplifying the Galaxy Evolution Picture with Age Matching} \label{subsec:simplicity} The primary result of this paper is that the simple age matching model, in which galaxies and halos co-evolve, such that quenched galaxies reside in old halos, is able to predict a wide variety of observed SDSS galaxy statistics for quenched and star-forming galaxies. In our model, there are (1) no fine tuning or updates to the age matching model that proved successful at reproducing color-based SDSS measurements, (2) no distinction between central and satellite galaxies when assigning SFRs, and (3) no explicit modeling of post-accretion processes that are believed to stifle the star formation in satellite galaxies (e.g., strangulation, ram pressure stripping, etc.). Let us consider these points in turn. As discussed in \S~\ref{sec:intro}, the color of a galaxy is known to be strongly correlated with star formation activity. For a variety of reasons, though, the correspondence is not perfect. For instance, active galaxies can often be classified as red due to the presence of dust \citep{maller_etal09,masters_etal10}. Color correlates with long term mass accretion history in age matching because of the timescale ($\sim\Gyr$) to evolve from the blue to red sequence. On the other hand, the timescales relevant to, for example, H$\alpha$ indicators of SFR are significantly shorter than timescales impacting color (e.g., $\sim10-100\Myr$ for the lifetime of O and B stars). Therefore, it is plausible that employing present day SFR in the age matching model may not exhibit the same level of success as a model based on $g - r$ color. We have shown that this is not the case: the SFR predictions of our age matching implementation of CAM are equally successful as the color-based predictions from Papers I $\&$ II. Certainly the relatively tight scatter between $g-r$ and sSFR is partly responsible for this dual success. In a follow-up paper to the present work (Watson, Skibba $\&$ Hearin 2014, in prep), we will show that the shortcomings of using broadband color as a proxy for present day star formation activity have virtually no manifestation on the two-point function. This surprising result, interesting in its own right, provides further insight into the reason that our model is able predict both $g-r$ color and SFR without modification. The explanation for this is simple: a star-forming galaxy appears red when our line of sight to the galaxy lies in the plane of its disk; for a {\em pair} of galaxies separated by $r\gtrsim100\kpc,$ the probability of this occurrence is essentially independent due to the very weak correlation between galaxy alignments \citep{zhang_etal13}. Point (2) highlights the simplicity of age matching, as well as what drives the satellite quenching predictions of the model. Consider the implications of the left panel of Fig.~\ref{fig:age_cen_sat_quenching}. We use our mock catalog to show the average formation epoch of centrals (mock galaxies residing in host halos) in comparison to satellites (subhalos). As in Papers I $\&$ II, we use the \citet{wechsler02} concentration-based definition for the formation epoch of a halo. In age matching, despite there being no distinction between central and satellite galaxies when assigning a SFR, satellite galaxies are more quenched than their central galaxy counterparts at fixed stellar mass simply because subhalos form earlier than host halos. The empirical justification for this cornerstone of age matching is illustrated in the right panel of Fig.~\ref{fig:age_cen_sat_quenching}, where we show the difference in the mean SFR of satellite and central galaxies in bins of fixed stellar mass. For this figure, we now use the group-finder to identify centrals and satellites, permitting a direct comparison to observational data. SDSS measurements are shown as black, filled circles and our age matching model prediction is shown as the solid black line with a gray error band (all errors are Poisson). At fixed stellar mass, satellites are more quenched than centrals in both the data and the model, and the observed quenching difference is quantitatively consistent with the difference implied by the relative formation times of host halos and subhalos. This observation is closely connected to point (3). Age matching does not require any explicit modeling of post-infall effects on satellite galaxy quenching. The virial radius $\Rvir$ of host halos only enters into our model through the definition of $\zacc,$ the epoch when a halo accretes onto a larger halo, thus becoming a subhalo. However, recall that in age matching, SFR is determined by $\zstarve$, the redshift in a (sub)halo's MAH when it is deemed to be starved of the cold gas supply needed to continue fueling star formation. Formally, $\zstarve\equiv \mathrm{Max}\left\{\zacc,\zchar,\zform\right\}$, and as we showed in Paper II, the epoch $\zacc$ has an essentially negligible impact on $\zstarve$ at all stellar masses, a result which also holds true in the present work. {\em Thus in our model, $\Rvir$ does not mark a special transition region in the evolution of a satellite,} and yet we accurately predict the radial profiles of quenched and star-forming galaxies both inside and well beyond the group radius, as well as the so-called ``excess quenched fraction'' of satellites (right panel of Fig.~\ref{fig:age_cen_sat_quenching}). \begin{figure*} \begin{center} \includegraphics[width=1.\textwidth]{./FIGS/age.stellarmass.cens.sat.specific.quenching.eps} \caption{{\bf Left Panel:} Formation epoch of central (blue curve) and satellite (red curve) galaxies as a function of stellar mass. Satellites in our model are more quenched than central galaxies of the same stellar mass simply because subhalos form earlier than host halos. This fact about structure formation in CDM is what drives satellite quenching in the age matching model. {\bf Right Panel:} The difference between the average SFR of satellite and central galaxies as a function of stellar mass for SDSS (filled, black circles) and our age matching prediction (black solid line). Poisson errors are shown for both data and the model. At fixed stellar mass satellites have lower SFRs than their central galaxy counterparts.} \label{fig:age_cen_sat_quenching} \end{center} \end{figure*} This qualitatively distinguishes age matching from conventional semi-analytic and empirical models of satellite quenching. We note, however, that the lack of explicit appearance of $\Rvir$ in our model does {\em not} imply that post-accretion processes are necessarily irrelevant to satellite quenching, since there is a significant correlation between the time a subhalo accretes onto a larger halo and the time the subhalo formed (see Fig.~6 of Paper II), rather that the overallin fluence of post-accretion physics has been overstated in the literature. Nonetheless, we will show in a pair of companion papers to this one that recent observations of SFR trends in the low-redshift universe {\em do} favor a scenario in which quenching is impacted by physical processes that operate on scales far larger than $\Rvir,$ {\em even for central galaxies}, as discussed in the following section. \subsection{Discriminating between Competing Quenching Models} \label{subsec:conformity} As discussed in detail in \citet*{zentner_etal14}, the success of age matching has exposed fundamental degeneracies in traditional approaches to galaxy-halo modeling such as the Halo Occupation Distribution \citep[HOD, e.g.,][]{seljak00,cooray02,berlind02,berlind03,zheng05,skibba_sheth09,watson_etal10,watson_etal12a} and Conditional Luminosity Function \citep[CLF, e.g.,][]{yang03,vdBosch13}. While it is true that age matching is based on $\vmax$ to set the stellar mass or luminosity content of halos, the assignment of the additional galaxy properties of color or SFR is based on halo assembly history. Conversely, HOD modeling beyond just stellar mass- or luminosity - dependent clustering, i.e. the color or SFR dependence \citep[e.g.,][]{zehavi05a,zehavi11,skibba09a,tinker_etal13,guo_SDSS14}, is still exclusively governed by $\Mvir$ and no other halo property. And yet, both classes of models give very good descriptions of a wide variety of measurements of the galaxy distribution. These considerations apply equally well to degeneracies with other common models of galaxy evolution. Indeed, both HODs and CLFs enjoy comparable levels of qualitative successes in reproducing observed statistics such as those presented in this work and the previous age matching papers. There is thus some legitimate cause for concern that conventional statistics describing the galaxy distribution are inadequate to conclusively discriminate between competing models. One particularly interesting measurement is that of {\em galactic conformity}, a feature in the galaxy distribution first discovered by \citet{weinmann06b}. In their analysis of an SDSS galaxy group catalog, \citet{weinmann06b} showed that in group systems of the same halo mass, satellites in groups with a red central tend to be redder than satellites in groups with a blue central. In another recently reported detection of conformity, \citet{kauffmann_etal13} showed that in an SDSS sample of central galaxies of the same stellar mass, the environment surrounding quenched central galaxies exhibits, on average, an attenuated SFR relative to the environment around star-forming centrals, a correlation that persists out to scales of $R\sim5\Mpc,$ far outside the virial radius of the host halo of the centrals. As we show in a recent paper \citep*{APH_DFW_vdB}, these closely related signals are formally distinct in the following sense: the \citet{weinmann06b} notion of conformity pertains to SFR correlations between central and satellite galaxies in the same dark matter halo, while the larger scale \citet{kauffmann_etal13} signal pertains to SFR correlations in distinct halos. We contend that no galaxy evolution model in which central galaxy SFR is exclusively determined by halo mass $\Mvir$ (and subsequently virial radius $\Rvir$) can account for either signal, as there would be no mechanism by which such correlations could arise. However, in age matching, galaxies in the same environment evolve from collapsed peaks of the same region of the initial cosmic density field. Thus the known correlation between the formation times of nearby halos \citep[e.g.,][]{sheth_tormen04,wechsler06} naturally gives rise to correlated stellar mass assembly histories of nearby galaxies. In \citet*{APH_DFW_vdB} we demonstrate that age matching predicts galactic conformity of both the \citet{weinmann06b} and \citet{kauffmann_etal13} varieties, with no modifications to the model presented in this work. Since it was shown in \citet{kauffmann_etal13} that the \citet{guo_etal11} semi-analytic model (SAM) does not predict conformity, this signal appears to be a promising testbed for the further development of galaxy evolution models.\footnote[5]{For example, the ``pre-heating'' of gas in the inter-galactic medium implemented in the SAM recently introduced in \citet{lu_etal14} is a promising mechanism by which conformity may arise, as discussed in \citet{kauffmann_etal13}.} \subsection{Future Directions} \label{subsec:future} This trilogy of age matching papers has revealed that there is a surprisingly simple relationship between the star formation activity of a galaxy and the assembly history of its dark matter halo. However, there are two clear paths to challenging the age matching picture of galaxy and halo co-evolution. First, our model has only been tested for central and satellite galaxies of $\log_{10}(\Mstar)>9.8$. However, using observations of classical dwarf galaxies in SDSS, \citet{geha_etal12} discovered that there is an apparent stellar mass threshold of $\log_{10}(\Mstar)=9.0,$ below which quenched galaxies do not exist in the field. In a recent study of this SDSS dwarf sample, \citet{wheeler_etal14} demonstrated that the so-called ``quenching timescale'' after a satellite first crosses the virial radius $\Rvir$ of its host halo must be implausibly long ($\gtrsim 9 \mathrm{Gyr}$) to produce the trends reported in \citet{geha_etal12}. These results are intriguingly in keeping with the notion supported by age matching that the role of $\Rvir$ has been over-estimated in the literature. In future work, we aim to apply the CAM modeling technique to dwarf galaxy samples to investigate whether the SFR trends exhibited by galaxies in this mass range are also reflected in a simple way by the evolutionary history of dark matter halos. The second consideration will be confronting age matching with observations at higher redshift. We will soon take this crucial next step thanks to high-completeness data sets such as PRIMUS \citep{Primus_Coil}, GAMA \citep{GAMA_Driver} and VIPERS \citep{VIPERS_Guzzo}. Ultimately, the power of this class of semi-empirical models is their ability to be used as ``training sets'' to help inform more complicated models of galaxy formation that include prescriptions for physical processes \emph{a priori}. Specifically, SAMs and hydrodynamic simulations can draw insight from age matching in order to improve their input physical recipes. In this paper we have explored the hypothesis that the star formation rates (SFRs) of galaxies can be determined via the ansatz that galaxies co-evolve with their dark matter halos. Specifically, we have studied the age matching formalism introduced in \citet{HW13a}, whose central tenet is that red/quenched galaxies reside in old halos, and conversely for blue/star-forming galaxies. This simple formalism has been proven to be remarkably powerful, yielding accurate predictions of SDSS color-dependent clustering and galaxy-galaxy lensing, as well as a variety of galaxy group-based statistics. In this work have confronted our age matching formalism with SFR-dependent, low-redshift galaxy statistics. Specifically, we have found the following principal results. \begin{itemize} \item{We present new measurements of SDSS clustering and galaxy-galaxy lensing as a function of stellar mass, split into distinct quenched and star-forming populations. The same age matching prescription introduced in Paper I and extended in Paper II adapts seamlessly to accurately predict these SFR-dependent observations, {\em without necessitating updates to the model or fine-tuning/fitting of parameters}.} \item{Age matching predictions are in excellent agreement with the observed radial distribution of star-forming and quenched satellite galaxies within and around galaxy group, rich group, and cluster environments, a success that extends significantly beyond the group radius.} \item{We demonstrate the lack of halo mass-dependence in the slope of the radial quenched fraction of satellites, finding an $\sim \mathrm{r}^{-.15}$ gradient {\em independent of environment}.} \item{We make our mock galaxy catalog publicly available at {\tt http://logrus.uchicago.edu/$\sim$aphearin}.} \end{itemize} These findings provide compelling evidence for the co-evolution of halos and galaxies, and are highly suggestive of the conclusion that the existing literature has over-estimated the role of post-accretion processes on attenuating star formation in satellite galaxies. We consider the myriad successes of our model to indicate that there does indeed exist a simple relation between cosmic star formation history of galaxies and the dark side of the universe.
14
3
1403.1578
In this paper, we test the age matching hypothesis that the star formation rate (SFR) of a galaxy of fixed stellar mass is determined by its dark matter halo formation history, e.g. more quiescent galaxies reside in older haloes. We present new Sloan Digital Sky Survey measurements of the galaxy two-point correlation function and galaxy-galaxy lensing as a function of stellar mass and SFR, separated into quenched and star-forming galaxy samples to test this simple model. We find that our age matching model is in excellent agreement with these new measurements. We also find that our model is able to predict: (1) the relative SFRs of central and satellite galaxies, (2) the SFR dependence of the radial distribution of satellite galaxy populations within galaxy groups, rich groups, and clusters and their surrounding larger scale environments, and (3) the interesting feature that the satellite quenched fraction as a function of projected radial distance from the central galaxy exhibits an ∼r<SUP>-.15</SUP> slope, independent of environment. These accurate predictions are intriguing given that we do not explicitly model satellite-specific processes after infall, and that in our model the virial radius does not mark a special transition region in the evolution of a satellite. The success of the model suggests that present-day galaxy SFR is strongly correlated with halo mass assembly history.
false
[ "satellite galaxy populations", "galaxy groups", "halo mass assembly history", "older haloes", "halo mass assembly", "galaxy-galaxy", "central and satellite galaxies", "present-day galaxy SFR", "projected radial distance", "fixed stellar mass", "stellar mass", "its dark matter halo formation history", "SFR", "environment", "the central galaxy", "rich groups", "new Sloan Digital Sky Survey measurements", "quenched and star-forming galaxy samples", "a galaxy", "satellite-specific processes" ]
11.817473
6.099261
-1
437678
[ "Varenius, E.", "Conway, J. E.", "Martí-Vidal, I.", "Aalto, S.", "Beswick, R.", "Costagliola, F.", "Klöckner, H. -R." ]
2014A&A...566A..15V
[ "The radio core structure of the luminous infrared galaxy NGC 4418. A young clustered starburst revealed?" ]
31
[ "Department of Earth and Space SciencesChalmers University of Technology, Onsala Space Observatory, 439 92, Onsala, Sweden", "Department of Earth and Space SciencesChalmers University of Technology, Onsala Space Observatory, 439 92, Onsala, Sweden", "Department of Earth and Space SciencesChalmers University of Technology, Onsala Space Observatory, 439 92, Onsala, Sweden", "Department of Earth and Space SciencesChalmers University of Technology, Onsala Space Observatory, 439 92, Onsala, Sweden", "Jodrell Bank Centre for Astrophysics, Alan Turing Building, School of Physics and Astronomy, The University of Manchester, Manchester, M13 9PL, UK", "Instituto de Astrofísica de Andalucía, Glorieta de la Astronomá, s/n, 18008, Granada, Spain", "Max-Planck-Institut für Radioastronomie, auf dem Hügel 69, 53121, Bonn, Germany" ]
[ "2015A&A...581L...5G", "2015A&A...582A..91C", "2015ApJ...800...69G", "2015MNRAS.448.3325J", "2015MNRAS.449.2598R", "2017A&A...607A..43V", "2017MNRAS.465.3943P", "2018A&A...609A..75F", "2018A&A...620A.112B", "2018ApJ...860...73E", "2018MNRAS.476.3478B", "2018arXiv181208953O", "2019ApJ...871..191O", "2019ApJ...882..153G", "2020ApJ...892...29F", "2020ApJ...904..200Y", "2020MNRAS.494.1784A", "2021A&A...648A..14R", "2021A&A...649A.105F", "2021ApJ...923..206S", "2021ApJ...923..240S", "2021MNRAS.502.3021R", "2022A&A...658A...3S", "2022ApJ...936...58Z", "2022ApJ...940...52S", "2023A&A...679A.115P", "2023ApJ...951...87O", "2023arXiv230911603J", "2024A&A...683A..27W", "2024ApJS..271...40J", "2024arXiv240612126G" ]
[ "astronomy" ]
7
[ "galaxies: Seyfert", "galaxies: star formation", "galaxies: individual: NGC 4418", "Astrophysics - Astrophysics of Galaxies", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
[ "1986MNRAS.218P..19R", "1988ApJ...326L..45A", "1989ApJ...347..727A", "1990ApJS...73..359C", "1991ApJ...378...65C", "1992ARA&A..30..575C", "1995PhDT.......238B", "1996AAS...18910905G", "1997ASPC..125...77S", "2001A&A...365L.353S", "2001ApJ...553..968J", "2001ApJ...554..803Y", "2001isra.book.....T", "2003AJ....125.2341E", "2003ASSL..285..109G", "2003ApJ...586..794B", "2003ApJ...593..733R", "2003JKAS...36..167S", "2003MNRAS.344L..59M", "2004AJ....127..239G", "2005ApJ...630..167T", "2006A&A...449..559B", "2006ASPC..351..497K", "2007A&A...461..471O", "2007ASPC..376..127M", "2009ApJ...691..946M", "2009tra..book.....W", "2011ApJ...740...95B", "2012A&A...539A.134B", "2012A&A...541A...4G", "2012A&A...545A..65P", "2013A&A...556A..66C", "2013ApJ...764...42S" ]
[ "10.1051/0004-6361/201323303", "10.48550/arXiv.1403.3782" ]
1403
1403.3782_arXiv.txt
The galaxy \object{NGC\,4418} (IRAS12243-0036) is a luminous ($L_{FIR}>10^{11}L_\sun$) infrared galaxy (LIRG) with an infrared flux density more than 5 times larger than expected from the linear radio to far-infrared relation. \cite{yun} find only ten such objects in a sample of 1809 galaxies, which makes NGC\,4418 a very unusual object. The galaxy has one of the deepest mid-IR silicate absorption features ever detected towards an external galaxy indicating a very deeply obscured nucleus. The deep absorption and high IR brightness of NGC\,4418 was first noted by \cite{roche1986} who proposed that this galaxy hides either an active galactic nucleus (AGN) or a compact nuclear starburst. Very faint H$\alpha$ emission has been detected, but the absence of NII, OI, and SII emission makes it difficult to classify the galaxy as an AGN or starburst based on its optical spectrum \citep{armus1989}. Furthermore, near-IR and mid-IR observations are consistent with both hypotheses \citep{evans2003}. Millimetre and submm observations by \cite{sakamoto2013} and \cite{francescolatest} reveal the presence of a highly compact (<0."1) high surface brightness continuum source suggesting that the bulk of the galaxy FIR emission emerges from a source less than 20\,pc in diameter. The extreme inferred H$_2$ column density $>10^{25}\text{cm}^{-2}$ towards this nucleus \citep{gonzales2012} makes it extremely difficult to determine the nature of the buried source. Despite extensive studies of NGC\,4418, the nature of its central power source is not clear. High-resolution radio very long baseline interferometry (VLBI) observations provide a possible way to distinguish between AGNs and starbursts. If the source of emission is an AGN this may produce high brightness temperature compact radio components. If the source of emission is a young starburst this would instead produce multiple supernovae (SNe) and supernova remnants (SNRs). A mix between these two scenarios, i.e. AGN and SNe/SNRs, is also possible. Radio VLBI observations with (sub)milli arcsecond resolution have proven to be a valuable tool to distinguish between these scenarios by directly probing the central regions (U)LIRGs, e.g. for Arp\,220 \citep{batejat2011} and Arp\,299 \citep{Arp299Paper}. In this paper we report on the analysis of archival NGC\,4418 data from the EVN and MERLIN interferometers which for the first time reveal eight discrete compact radio features within its nucleus. In Sect.\,\ref{sec:cal} we summarise the data used and the calibration procedures applied. In Sect.\,\ref{sec:imaging} we describe the imaging process and discuss the image fidelity. In Sect.\,\ref{sec:results} we present the results of the imaging and simple modelling. In Sect.\,\ref{sec:discussion} we briefly discuss the hypotheses of AGN/starburst in relation to our results. Finally, in Sect.\,\ref{sec:conclusions} we summarise our conclusions. In this paper we assume a distance of 34\,Mpc to NGC\,4418 at which an angular size of 1 mas corresponds to 0.165\,pc \citep{francescolatest}.
\label{sec:conclusions} Eight compact features have been detected in the nucleus of NGC\,4418 at 5.0\,GHz. The complex morphology and inverted spectrum of the compact features can be seen as evidence against the hypothesis that an active galactic nucleus alone is powering the nucleus of NGC\,4418. This indicates a significant contribution from star formation. The compact features can be super star clusters with intense star formation. We note, however, that the surface brightness of the compact features is close to the limit of what can be produced by well-mixed thermal/non-thermal emission from any surface density of star-formation. This could be due to an AGN responsible for some of the radio emission while the rest is due to star formation. Unfortunately, the current data does not allow us to clearly separate an AGN feature from features due to star formation. New multi frequency radio VLBI observations with the current capabilities of EVN and eMERLIN are planned in 2014. We hope that these will further constrain the nature of the compact features in the nucleus of NGC\,4418. \vspace{1cm}
14
3
1403.3782
Context. The galaxy NGC 4418 contains one of the most compact obscured nuclei within a luminous infrared galaxy (LIRG) in the nearby Universe. This nucleus contains a rich molecular gas environment and an unusually high ratio of infrared-to-radio luminosity (q-factor). The compact nucleus is powered by either a compact starburst or an active galactic nucleus (AGN). <BR /> Aims: The aim of this study is to constrain the nature of the nuclear region (starburst or AGN) within NGC 4418 via very-high-resolution radio imaging. <BR /> Methods: Archival data from radio observations using the European Very Long Baseline Interferometry Network (EVN) and Multi-Element Radio Linked Interferometer Network (MERLIN) interferometers are imaged. Sizes and flux densities are obtained by fitting Gaussian intensity distributions to the image. The average spectral index of the compact radio emission is estimated from measurements at 1.4 GHz and 5.0 GHz. <BR /> Results: The nuclear structure of NGC 4418 visible with EVN and MERLIN consists of eight compact (&lt;49 mas i.e. &lt;8 pc) features spread within a region of 250 mas, i.e. 41 pc. We derive an inverted spectral index α ≥ 0.7 (S<SUB>ν</SUB> ∝ ν<SUP>α</SUP>) for the compact radio emission. <BR /> Conclusions: Brightness temperatures &gt;10<SUP>4.8</SUP> K indicate that these compact features cannot be HII-regions. The complex morphology and inverted spectrum of the eight detected compact features is evidence against the hypothesis that an AGN alone is powering the nucleus of NGC 4418. The compact features could be super star clusters with intense star formation, and their associated free-free absorption could then naturally explain both their inverted radio spectrum and the low radio-to-IR ratio of the nucleus. The required star formation area density is extreme, however, and close to the limit of what can be observed in a well-mixed thermal/non-thermal plasma produced by star formation, and is also close to the limit of what can be physically sustained.
false
[ "radio observations", "radio", "Multi-Element Radio Linked Interferometer Network", "star formation", "intense star formation", "q-factor", "nuclei", "1.4 GHz", "5.0 GHz", "super star clusters", "the compact radio emission", "NGC", "inverted spectrum", "fitting Gaussian intensity distributions", "MERLIN", "AGN", "the European Very Long Baseline Interferometry Network", "The compact nucleus", "EVN", "Universe" ]
14.13955
7.943445
-1
745694
[ "Barnes, J. E.", "Wood, Kenneth", "Hill, Alex S.", "Haffner, L. M." ]
2014MNRAS.440.3027B
[ "Photoionization and heating of a supernova-driven turbulent interstellar medium" ]
24
[ "School of Physics &amp; Astronomy, University of St Andrews, North Haugh, St Andrews, Fife KY16 9SS, Scotland", "School of Physics &amp; Astronomy, University of St Andrews, North Haugh, St Andrews, Fife KY16 9SS, Scotland", "CSIRO Astronomy &amp; Space Science, Marsfield, NSW 1710, Australia", "Department of Astronomy, University of Wisconsin Madison, 475 North Charter Street, Madison, WI 53706, USA; Space Science Institute, 4750 Walnut Street, Suite 205, Boulder, CO 80301, USA" ]
[ "2014ApJ...787..106H", "2015HEDP...17..146C", "2015MNRAS.447..559B", "2016MNRAS.457.1257H", "2017A&A...599A.141J", "2017ApJ...838...43K", "2017ApJ...844..155T", "2017MNRAS.467.1067D", "2017suex.book.....B", "2018A&C....23...40V", "2018ApJ...862...55H", "2018MNRAS.476.4032V", "2018MNRAS.479.3042G", "2019A&A...622A.115W", "2019ApJ...882...84L", "2019HEDP...3300710L", "2019MNRAS.488.1977V", "2019MNRAS.489.5530R", "2020ApJ...897..143K", "2021ApJ...907...22T", "2021ApJ...917...72L", "2021ApJ...922..131T", "2022RASTI...1....3L", "2024arXiv240501637D" ]
[ "astronomy" ]
3
[ "ISM: General", "galaxies: ISM", "Astrophysics - Astrophysics of Galaxies" ]
[ "1963AuJPh..16....1H", "1977ApJ...218..148M", "1979IAUS...83..103U", "1982ApJ...263..777G", "1989MNRAS.239..605K", "1990ARA&A..28..215D", "1990ApJ...349L..17R", "1991PASP..103.1182A", "1993ApJ...417..579M", "1995ApJ...443..152W", "1995ApJ...443..209A", "1996ApJ...460..914V", "1996ApJ...462..563N", "1997ApJ...476..166W", "1997ApJ...477..196E", "1997ApJ...485..182M", "1997ApJ...490..143B", "1998ApJ...494L..99R", "1998ApJ...501..137R", "1999ApJ...523..223H", "1999ApJ...525L..21R", "2000A&A...363....9Z", "2000ApJ...544..347M", "2000ApJS..131..273F", "2000MNRAS.315..479D", "2001A&A...375..161P", "2001AJ....121.2737M", "2002A&A...386..801Z", "2002ApJ...565.1060H", "2002ApJ...572..823O", "2003ApJ...599.1333S", "2003ApJS..149..405H", "2004ApJ...611..338S", "2004MNRAS.348.1337W", "2004MNRAS.353.1126W", "2004RvMP...76..125M", "2005ARA&A..43..337C", "2005ApJ...630..925M", "2005ApJ...633..295W", "2006A&A...449..509H", "2006ApJ...652..401M", "2006ApJ...653.1266J", "2008ApJ...686..363H", "2008PASA...25..184G", "2008RPPh...71d6901K", "2009AJ....138.1577D", "2009ApJ...700.1299J", "2009ApJ...702.1472S", "2009ApJ...704..137J", "2009RvMP...81..969H", "2010ApJ...710..853C", "2010ApJ...721.1397W", "2012ApJ...749..145B", "2012ApJ...750..104H", "2012ApJ...760...91P", "2012ApJ...761..189H", "2012MNRAS.423L.107D", "2013ApJ...767...87W", "2013ApJ...771..132B", "2013ApJ...775..109K" ]
[ "10.1093/mnras/stu521", "10.48550/arXiv.1403.3261" ]
1403
1403.3261_arXiv.txt
The interstellar medium (ISM) is a vital component of the cycle of star formation and the evolution of galaxies. The composition and dynamics of the ISM determine the formation of new stars in the Galaxy, while stars provide feedback through ionisation, outflows and supernovae \citep{MacLow2004}. In this paper we study the formation of widespread ionised gas as observed primarily through \ha in the Milky Way and other galaxies. This gas (reviewed by \citealt{Haffner2009}), commonly referred to as the Diffuse Ionised Gas (DIG) or Warm Ionised Medium (WIM), is low density ($\la 0.1$~cm$^{-3}$), warm ($\sim8000$~K), consists of regions of nearly fully ionised hydrogen (\citealt{Hausen2002}) and has a scale height of $1-1.5$ kpc near the sun (\citealt{Savage2009}, \citealt{Haffner1999}, \citealt{Gaensler2008}). The most likely sources of the ionisation of the DIG in the Galaxy are O stars \citep{Reynolds1990}. Photoionisation simulations of a smooth ISM are able to reproduce some of the observed properties of the DIG (e.g., \citealt{Wood2004b}, \citealt{Miller1993}). However, to allow ionising photons from midplane OB stars to propagate to large distances, these models require the vertical column density of hydrogen to be lower than that inferred from HI 21 cm observations of the Galaxy. Photoionisation simulations of a clumpy or fractal density structure show that the introduction of lower density paths in a three dimensional (3D) ISM allow photons to travel to large heights above the midplane (see for example figure 16 of \citealt{Haffner2009}). The most likely source of such clumping is turbulence (e.g., \citealt{Armstrong1995}, \citealt{Hill2008}, \citealt{Chepurnov2010}, \citealt{Burkhart2012}) which could be driven by supernovae (e.g., \citealt{MacLow2004}, \citealt{Avillez2000}, \citealt{Armstrong1995}). \cite{Wood2010} demonstrated that in a 3D supernova-driven, turbulent medium, ionising photons are indeed able to propagate to large distances and produce the DIG. However, the width of the distribution of emission measure in these simulations is wider than observed in the Galaxy. The discrepancy appears to be due to too wide a variation of density with height, requiring a mechanism to smooth out the density variations in the dynamical simulations. One possible smoothing mechanism is pressure from magnetic fields. In this paper we extend the work of \cite{Wood2010} to study photoionisation of a supernova-driven, turbulent magnetised ISM, using the 3D density structures from the MHD simulations of \citet{Hill2012}. Our simulations naturally produce a vertically extended ionised component and a compact neutral layer of gas, in qualitative agreement with observations. However the \ha intensity from the ionised layer has a smaller scale height than observed in the Galaxy. To better reproduce \ha observations in the Galaxy we have created models of a 3D fractal ISM which provide estimates for the density structure and column densities for future MHD simulations. We also investigate the distance travelled by ionising photons in the ISM and find that the majority of photons travel only short distances and only a small fraction need to travel kiloparsec distances to ionise the DIG. The outline of the paper is as follows: a description of observations of the DIG is given in section \ref{WHAM}. The MHD and Monte Carlo photoionisation codes are briefly described in section \ref{models}, the results of our simulations are presented in section \ref{results} and are compared with observations of the DIG in the Galaxy. In section \ref{analytic} we describe the results of photoionisation models of a 3D fractal ISM. In section \ref{results:paths} we investigate how far photons are able to travel through the ISM to create the DIG. Finally our conclusions are presented in section \ref{conclusions}.
\label{conclusions} We have produced photoionisation simulations of the DIG in an environment similar to that in the outer disc of a spiral galaxy and compared these to observations of the Perseus Arm and an inter-arm region in the solar neighbourhood. We summarise our main conclusions here: \begin{enumerate} \item The photoionisation of 3D density structures from MHD simulations naturally produces widespread diffuse ionised gas with a density scale height larger than that of neutral hydrogen. However the density grids from the MHD simulations we have used here have a low scale height resulting in smaller \ha intensity scale height than observed in the Perseus Arm. \item We find that with the addition of a heating term, such as heating by cosmic rays or dissipation of turbulence, our simulations are able to reproduce general trends in \n optical emission line ratios in the DIG. \item A fractal density structure for the ISM with higher density at large $|z|$ than in the MHD simulations better reproduces WHAM observations of the Perseus arm. This will provide a guide for the required density structure for future MHD simulations of the DIG. \item Finally our simulations demonstrate that ionising photons are able to travel many kiloparsecs to ionise the DIG at large heights above the midplane, with photons travelling through low density ``bubbles'' close to the midplane and through low density diffuse gas at large heights. An important next step in modelling the ionisation of the DIG will be to include photoionisation as a dynamical process in MHD simulations of the gas (e.g., \citealt{deavillez2012}). This may increase the density of gas at large heights since photoionisation will increase the temperature of the gas and allow the gas to expand to a larger height, thereby sustaining higher densities above the midplane of the Galaxy that are demanded by the WHAM observations. \end{enumerate}
14
3
1403.3261
The diffuse ionized gas (DIG) in galaxies traces photoionization feedback from massive stars. Through three-dimensional photoionization simulations, we study the propagation of ionizing photons, photoionization heating and the resulting distribution of ionized and neutral gas within snapshots of magnetohydrodynamic simulations of a supernova-driven turbulent interstellar medium. We also investigate the impact of non-photoionization heating on observed optical emission line ratios. Inclusion of a heating term which scales less steeply with electron density than photoionization is required to produce diagnostic emission line ratios similar to those observed with the Wisconsin Hα Mapper. Once such heating terms have been included, we are also able to produce temperatures similar to those inferred from observations of the DIG, with temperatures increasing to above 15 000 K at heights |z| ≳ 1 kpc. We find that ionizing photons travel through low-density regions close to the mid-plane of the simulations, while travelling through diffuse low-density regions at large heights. The majority of photons travel small distances (≲100 pc); however some travel kiloparsecs and ionize the DIG.
false
[ "observed optical emission line ratios", "diagnostic emission line ratios", "photoionization heating", "photoionization feedback", "massive stars", "ionizing photons", "large heights", "magnetohydrodynamic simulations", "photoionization", "heights", "electron density", "non-photoionization heating", "such heating terms", "DIG", "temperatures", "≲100 pc", "the Wisconsin Hα Mapper", "diffuse low-density regions", "photons", "galaxies" ]
11.976143
8.84871
176
734949
[ "Will, Clifford M." ]
2014LRR....17....4W
[ "The Confrontation between General Relativity and Experiment" ]
2,165
[ "Department of Physics, University of Florida, 32611, Gainesville, FL, USA" ]
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[ "astronomy", "physics" ]
69
[ "Gravitational radiation", "Theories of gravity", "Post-Newtonian limit", "Tests of relativistic gravity", "General Relativity and Quantum Cosmology", "Astrophysics - Cosmology and Nongalactic Astrophysics", "High Energy Physics - Theory" ]
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"2014PhRvD..89j2002H", "2014PhRvD..90b2004M", "2014PhRvD..90d4025S", "2014PhRvL.112p1101Y", "2014SSRv..183..277R", "2014arXiv1402.5594W", "2014grav.book.....P", "2015AmJPh..83...39D" ]
[ "10.12942/lrr-2014-4", "10.48550/arXiv.1403.7377" ]
1403
1403.7377.txt
\label{S1} When general relativity was born 100 years ago, experimental confirmation was almost a side issue. Admittedly, Einstein did calculate observable effects of general relativity, such as the perihelion advance of Mercury, which he knew to be an unsolved problem, and the deflection of light, which was subsequently verified. But compared to the inner consistency and elegance of the theory, he regarded such empirical questions as almost secondary. He famously stated that if the measurements of light deflection disagreed with the theory he would ``feel sorry for the dear Lord, for the theory {\em is} correct!''. By contrast, today experimental gravitation is a major component of the field, characterized by continuing efforts to test the theory's predictions, both in the solar system and in the astronomical world, to detect gravitational waves from astronomical sources, and to search for possible gravitational imprints of phenomena originating in the quantum, high-energy or cosmological realms. The modern history of experimental relativity can be divided roughly into four periods: Genesis, Hibernation, a Golden Era, and the Quest for Strong Gravity. The Genesis (1887\,--\,1919) comprises the period of the two great experiments which were the foundation of relativistic physics -- the Michelson--Morley experiment and the E\"otv\"os experiment -- and the two immediate confirmations of general relativity -- the deflection of light and the perihelion advance of Mercury. Following this was a period of Hibernation (1920\,--\,1960) during which theoretical work temporarily outstripped technology and experimental possibilities, and, as a consequence, the field stagnated and was relegated to the backwaters of physics and astronomy. But beginning around 1960, astronomical discoveries (quasars, pulsars, cosmic background radiation) and new experiments pushed general relativity to the forefront. Experimental gravitation experienced a Golden Era (1960\,--\,1980) during which a systematic, world-wide effort took place to understand the observable predictions of general relativity, to compare and contrast them with the predictions of alternative theories of gravity, and to perform new experiments to test them. New technologies -- atomic clocks, radar and laser ranging, space probes, cryogenic capabilities, to mention only a few -- played a central role in this golden era. The period began with an experiment to confirm the gravitational frequency shift of light (1960) and ended with the reported decrease in the orbital period of the Hulse-Taylor binary pulsar at a rate consistent with the general relativistic prediction of gravitational-wave energy loss (1979). The results all supported general relativity, and most alternative theories of gravity fell by the wayside (for a popular review, see~\cite{WER}). Since that time, the field has entered what might be termed a Quest for Strong Gravity. Much like modern art, the term ``strong'' means different things to different people. To one steeped in general relativity, the principal figure of merit that distinguishes strong from weak gravity is the quantity $\epsilon \sim GM/Rc^2$, where $G$ is the Newtonian gravitational constant, $M$ is the characteristic mass scale of the phenomenon, $R$ is the characteristic distance scale, and $c$ is the speed of light. Near the event horizon of a non-rotating black hole, or for the expanding observable universe, $\epsilon \sim 1$; for neutron stars, $\epsilon \sim 0.2$. These are the regimes of strong gravity. For the solar system, $\epsilon < 10^{-5}$; this is the regime of weak gravity. An alternative view of ``strong'' gravity comes from the world of particle physics. Here the figure of merit is $GM/R^3c^2 \sim \ell^{-2}$, where the Riemann curvature of spacetime associated with the phenomenon, represented by the left-hand-side, is comparable to the inverse square of a favorite length scale $\ell$. If $\ell$ is the Planck length, this would correspond to the regime where one expects conventional quantum gravity effects to come into play. Another possible scale for $\ell$ is the TeV scale associated with many models for unification of the forces, or models with extra spacetime dimensions. From this viewpoint, strong gravity is where the curvature is comparable to the inverse length squared. Weak gravity is where the curvature is much smaller than this. The universe at the Planck time is strong gravity. Just outside the event horizon of an astrophysical black hole is weak gravity. Considerations of the possibilities for new physics from either point of view have led to a wide range of questions that will motivate new tests of general relativity as we move into its second century: \begin{itemize} \item Are the black holes that are in evidence throughout the universe truly the black holes of general relativity? \item Do gravitational waves propagate with the speed of light and do they contain more than the two basic polarization states of general relativity? \item Does general relativity hold on cosmological distance scales? \item Is Lorentz invariance strictly valid, or could it be violated at some detectable level? \item Does the principle of equivalence break down at some level? \item Are there testable effects arising from the quantization of gravity? \end{itemize} In this update of our {\em Living Review} , we will summarize the current status of experiments, and attempt to chart the future of the subject. We will not provide complete references to early work done in this field but instead will refer the reader to selected recent papers and to the appropriate review articles and monographs, specifically to \textit{Theory and Experiment in Gravitational Physics}~\cite{tegp}, hereafter referred to as TEGP; references to TEGP will be by chapter or section, e.g., ``TEGP~8.9''. Additional reviews in this subject are~\cite{PhysRevD.86.010001,shapiro,2008ARNPS..58..207T}. The ``Resource Letter'' by the author \cite{2010AmJPh..78.1240W}, contains an annotated list of 100 key references for experimental gravity. \newpage %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%
\label{S5} General relativity has held up under extensive experimental scrutiny. The question then arises, why bother to continue to test it? One reason is that gravity is a fundamental interaction of nature, and as such requires the most solid empirical underpinning we can provide. Another is that all attempts to quantize gravity and to unify it with the other forces suggest that the standard general relativity of Einstein may not be the last word. Furthermore, the predictions of general relativity are fixed; the pure theory contains no adjustable constants so nothing can be changed. Thus every test of the theory is either a potentially deadly test or a possible probe for new physics. Although it is remarkable that this theory, born 100 years ago out of almost pure thought, has managed to survive every test, the possibility of finding a discrepancy will continue to drive experiments for years to come. These experiments will search for new physics beyond Einstein at many different scales: the large distance scales of the astrophysical, galactic, and cosmological realms; scales of very short distances or high energy; and scales related to strong or dynamical gravity. \newpage %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%
14
3
1403.7377
The status of experimental tests of general relativity and of theoretical frameworks for analyzing them is reviewed and updated. Einstein's equivalence principle (EEP) is well supported by experiments such as the Eötvös experiment, tests of local Lorentz invariance and clock experiments. Ongoing tests of EEP and of the inverse square law are searching for new interactions arising from unification or quantum gravity. Tests of general relativity at the post-Newtonian level have reached high precision, including the light deflection, the Shapiro time delay, the perihelion advance of Mercury, the Nordtvedt effect in lunar motion, and frame-dragging. Gravitational wave damping has been detected in an amount that agrees with general relativity to better than half a percent using the Hulse-Taylor binary pulsar, and a growing family of other binary pulsar systems is yielding new tests, especially of strong-field effects. Current and future tests of relativity will center on strong gravity and gravitational waves.
false
[ "new tests", "general relativity", "experiments", "other binary pulsar systems", "gravitational waves", "strong gravity", "lunar motion", "Tests", "tests", "experimental tests", "Ongoing tests", "new interactions", "relativity", "local Lorentz invariance and clock experiments", "high precision", "frame-dragging", "theoretical frameworks", "Nordtvedt", "Mercury", "strong-field effects" ]
9.28831
1.936677
-1
482857
[ "Steele, Matthew M.", "Zepf, Stephen E.", "Maccarone, Thomas J.", "Kundu, Arunav", "Rhode, Katherine L.", "Salzer, John J." ]
2014ApJ...785..147S
[ "Composition of an Emission Line System in Black Hole Host Globular Cluster RZ2109" ]
15
[ "Department of Physics and Astronomy, Michigan State University, East Lansing, MI 48824, USA; Department of Physics, Northern Michigan University, Marquette, MI 49855, USA;", "Department of Physics and Astronomy, Michigan State University, East Lansing, MI 48824, USA", "Department of Physics, Texas Tech University, Lubbock, TX 79409, USA", "Eureka Scientific, 2452 Delmer Street Suite 100, Oakland, CA 94602-3017, USA", "Department of Astronomy, Indiana University, Bloomington, IN 47405, USA", "Department of Astronomy, Indiana University, Bloomington, IN 47405, USA" ]
[ "2018ApJ...862..108D", "2018MNRAS.476.1889T", "2019ApJ...877...57Q", "2019ApJ...883...44W", "2019MNRAS.485.1694D", "2019MNRAS.489.4783D", "2021MNRAS.504.1545D", "2021MNRAS.507..330S", "2021PhRvD.104l4024Y", "2022MNRAS.509.2493K", "2022MNRAS.512.3284S", "2023ApJ...951...91C", "2023LRR....26....2A", "2023NewAR..9601672K", "2024MNRAS.529.1347D" ]
[ "astronomy" ]
9
[ "galaxies: individual: NGC 4472", "galaxies: star clusters: individual: NGC 4472", "globular clusters: general", "X-rays: binaries", "X-rays: galaxies: clusters", "Astrophysics - Astrophysics of Galaxies" ]
[ "1998PASP..110..761F", "2001MNRAS.322..231K", "2003ApJ...596..997S", "2004A&A...425..881L", "2005ApJ...627..767M", "2006MNRAS.366.1151S", "2007ApJ...669L..69Z", "2007Natur.445..183M", "2007astro.ph..3658P", "2008A&A...477L..17L", "2008ApJ...683L.139Z", "2008MNRAS.386.2075S", "2009ApJ...705L.168G", "2010AIPC.1314..143I", "2011ApJ...739...95S", "2012ApJ...759..126P" ]
[ "10.1088/0004-637X/785/2/147", "10.48550/arXiv.1403.2784" ]
1403
1403.2784_arXiv.txt
The globular cluster RZ2109 located in the galaxy NGC4472 is a known host of an accreting black hole system, first identified by \citet{Maccarone2007}. Along with the variable X-ray source indicative of the accreting black hole, RZ2109 has been observed to host a broad and luminous [O~III]$\lambda \lambda$4959,5007 emission line complex thought to correspond to an accretion-powered outflow driven from the cluster's black hole \citep{Zepf2007,Zepf2008,Steele2011}. The RZ2109 emission line structure is unusual in many respects. The lines are very broad with a width of 3200 km s$^{-1}$, about a factor of 100 larger than the escape velocity of the globular cluster in which it is located. The line velocity profiles also have a complex shape and appear to contain two distinct velocity structures \citep{Steele2011}. Moreover, [O~III]$\lambda \lambda$4959,5007 are the only emission lines apparent even in a fairly deep spectrum \citep{Zepf2008}. \citet{Gnedin2009} suggested that the lack of Balmer line emission, in particular, may be indicative of a hydrogen poor donor star such as a white dwarf. This possibility is significant because although BH-WD binaries are an expected end stage of stellar evolution, especially in globular clusters, none has yet been positively identified \citep{Ivanova2011}. A common diagnostic for the study of emission line regions is the \ovh emission line ratio. Previous work by \citet{Zepf2008} noted that \ovh ratio appeared to be at least 30 based on a low resolution optical spectrum of the RZ2109 cluster. For reference other common astrophysical [O~III] emission line production sites include active galactic nuclei with typical \ovh ratios of order unity \citep{Sarzi2006}, Milky Way planetary nebulae with mean line ratios of approximately 15 \citep{Stanghellini2003}, and a few extremely hydrogen poor planetary nebulae with line ratio of ~20 \citep{Mendez2005,Larsen2008}. Based on a comparison to these sources, it seems plausible that the RZ2109 emission line site would be hydrogen-depleted. On the other hand, the [O III] production sites and mechanisms can be very different among astrophysical object classes and it is not immediately obvious how relevant such comparisons are. In this work we investigate the composition of the RZ2109 emission line region in order to help constrain models of the black hole/donor star binary in the globular cluster. Accordingly, Section \ref{sec:obs} presents an analysis of higher signal-to-noise spectra of RZ2109 than have been previously used to explore the composition of the source. In Section \ref{sec:ar}, we begin by modeling the stellar component of the RZ2109 spectrum to remove the stellar contributions from the emission line measurements. The next section presents the line measurement results. Section \ref{sec:rtm} presents results from a suite of radiative models that we use to interpret the range of \ovh emission line ratios. The last section of the paper includes a discussion of the overall results of this analysis.
\label{sec:dc} From the EW measurements presented in Section \ref{sec:ew} it is clear that any H$\beta$ emission which may be present in the observed RZ2109 spectrum is extremely weak. The H$\beta$ emission is sufficiently weak to make interpretation of the resulting \ovh line ratios difficult. If taken at face value these ratios are among the largest detected from any astrophysical source \citep{Sarzi2006,Mendez2005}. In Table \ref{tab:lim} the confidence levels for the \ovh ratio are given based on the uncertainty in the H$\beta$ measurement. Equivalent width measurements below the level of the continuum (absorption features for example) are indicated with negative equivalent width value. Therefore a negative line ratio limit in Table \ref{tab:lim} indicates that the H$\beta$ measurement is consistent with an absorption line at the specified confidence level. When the full velocity width covered by the [O~III]$\lambda$5007 complex is considered along with the radiative transfer modeling presented in Section \ref{sec:rtm} it is clear that a solar composition gas is insufficient to produce the observed \ovh ratios. The measured ratio using 3200 \kms apertures is nearly a factor of three larger than the maximum ratios produced by the synthetic emission line models. At the $95\%$ confidence level for the 3200 \kms aperture and the $90\%$ confidence level for the 600 \kms aperture the uncertainty limits on the \ovh ratio approach the maximum synthetic \ovh values. These ratio limits are 35.7 at 3200 km~s$^{-1}$, and 33.0 at 600 \kms compared to the synthetic maximum of 34. As such it may be possible to produce the necessary \ovh ratio for either aperture given an emission line region model that falls in a very specific location is physical parameter space. It should be noted, however, that the gas masses required to produce the maximum synthetic \ovh values are well above what might be expected to be associated with an accreting black hole in a globular cluster. In order to produce the maximum synthetic \ovh ratios a total gas mass of order 100 M$_{\sun}$ is needed which would eliminate X-ray binaries, planetary nebulae, supernovae remnants, or any other stellar scale objects as the source of the emitting gas. To reach a gas mass of that size a significant portion of the gas would necessarily be contributed by the cluster's interstellar medium. However the hydrogen densities involved in producing the maximum synthetic \ovh ratios are of order 10$^4$--10$^5$ cm$^{-3}$, orders of magnitude above the expected densities of a cluster interstellar medium. Furthermore, the 10$^{-1}$ solar composition of an interstellar medium would produce lower \ovh ratios in comparison to the solar composition models considered above, making a large contribution of ISM material to the emission line region unlikely. More typical gas masses produce maximum ratios a factor of four times smaller than those observed in the 3200 \kms aperture and nearly twice that of the lower limit of the 600 \kms measurement. The RZ2109 emission line region therefore appears to be oxygen enriched relative to solar composition. The level of oxygen enrichment necessary to produce the observed \ovh is not readily apparent, as the emission line systems physical parameters are not well constrained by observation. However, it is possible to examine some interesting limits by calculating the affect of carbon and oxygen enrichment for two characteristic masses; 1.0 M$_{\sun}$ representing a stellar source of the emitting gas, and 0.1 M$_{\sun}$ representing the outflow from a CO white dwarf and black hole binary. A grid of cloudy models, as described in section 3.3, was run for each characteristic mass with carbon and oxygen enriched from solar composition to 100 times solar. From this calculation we find a minimum CO enrichment of 17.8 times solar is necessary for the 1.0 M$_{\sun}$ model to yield \ovh ratios equivalent to the measured values of the 3200 \kms aperture and 5.5 times solar to match the 600 \kms aperture. For the 0.1 M$_{\sun}$ model the gas must be enriched to 17.2 times solar for the 3200 \kms aperture and 5.6 times solar to match the 600 \kms aperture. We must emphasize that these are only lower limits on CO enrichment for two interesting astrophysical systems, and are not necessarily indicative of enrichment of the RZ2109 emission line source. \citet{Steele2011} present the argument that the complex line [O~III] velocity profile observed in the RZ2109 spectrum is consistent with emission originating from two discrete gas structures. The two velocity apertures presented in this work do not directly correspond to the two geometric components discussed by \citet{Steele2011} as emission from the two components are superimposed in velocity space. From a comparison of the measured \ovh line ratios and synthetic line ratio considered here it seems most likely that both the gas structures are comprised of oxygen enriched material. For the higher velocity structure which \citet{Steele2011} describes as a bipolar conical outflow this is consistent with the scenario of material being stripped from a CO white dwarf companion to the black hole, and driven to the observed velocity as an accretion-powered outflow. The added constraint of being oxygen enriched, does not, however place obvious constraints on the gas source for the lower velocity component. As evidenced by the calculations summarized in Figure \ref{fig:mvr}, the geometry of the emission region strongly influences the emission line ratios that it produces. As such the \ovh ratio alone is insufficient to fully constrain the gas composition of the RZ2109 emission line system, or to positively identify the particular stellar type of the X-ray binary's donor star. The WD donor star model is consistent with all the observations and calculations presented above. However we are not yet able to fully rule out other late stage or polluted stellar types with super-solar abundance atmospheres as the source for the observed outflow. With future observations of other spectral bands, including measurements of UV carbon lines CIV 1548 \& 1551 \AA\ and CIII] 1907 \& 1910 \AA, and more precise estimates of the emission line system's gas mass, it may yet be possible to place tighter constraints on the emission line region's composition and the X-ray binary source system.
14
3
1403.2784
We present an analysis of optical spectra from the globular cluster RZ2109 in NGC 4472, which hosts the first unambiguous globular cluster black hole. We use these spectra to determine the elemental composition of the emission line system associated with this source, and to constrain the age and metallicity of the host globular cluster. For the emission line system of RZ2109, our analysis indicates the [O III] λ5007 equivalent width is 33.82 ± 0.39 Å and the Hβ equivalent width is 0.32 ± 0.32 Å, producing a formal [O III] λ5007/Hβ emission line ratio of 106 for a 3200 km s<SUP>-1</SUP> measurement aperture covering the full velocity width of the [O III] λ5007 line. Within a narrower 600 km s<SUP>-1</SUP> aperture covering the highest luminosity velocity structure in the line complex, we find [O III] λ5007/Hβ = 62. The measured [O III] λ5007/Hβ ratios are significantly higher than can be produced in radiative models of the emission line region with solar composition, and the confidence interval limits exclude all but models which have gas masses much larger than those for a single star. Therefore, we conclude that the region from which the [O III] λ5007 emission originates is hydrogen-depleted relative to solar composition gas. This finding is consistent with emission from an accretion-powered outflow driven by a hydrogen-depleted donor star, such as a white dwarf, being accreted onto a black hole.
false
[ "the first unambiguous globular cluster black hole", "solar composition gas", "solar composition", "gas masses", "emission", "Hβ", "a black hole", "radiative models", "the emission line system", "the emission line region", "RZ2109", "the globular cluster RZ2109", "the Hβ equivalent width", "the host globular cluster", "NGC", "a single star", "first", "the full velocity width", "the highest luminosity velocity structure", "a formal [O III] λ5007/Hβ emission line ratio" ]
6.904605
7.068112
-1
437465
[ "Moehler, S.", "Dreizler, S.", "LeBlanc, F.", "Khalack, V.", "Michaud, G.", "Richer, J.", "Sweigart, A. V.", "Grundahl, F." ]
2014A&A...565A.100M
[ "Hot horizontal branch stars in NGC 288 - effects of diffusion and stratification on their atmospheric parameters" ]
13
[ "European Southern Observatory, Karl-Schwarzschild-Str. 2, 85748, Garching, Germany ; Institut für Theoretische Physik und Astrophysik, Olshausenstraße 40, 24118, Kiel, Germany", "Georg-August-Universität, Institut für Astrophysik, Friedrich-Hund-Platz 1, 37077, Göttingen, Germany", "Département de Physique et d'Astronomie, Université de Moncton, Moncton, New Brunswick, E1A 3E9, Canada", "Département de Physique et d'Astronomie, Université de Moncton, Moncton, New Brunswick, E1A 3E9, Canada", "Département de physique, Université de Montréal, Montréal, Québec, H3C 3J7, Canada", "Département de physique, Université de Montréal, Montréal, Québec, H3C 3J7, Canada", "NASA Goddard Space Flight Center, Exploration of the Universe Division, Code 667, Greenbelt, MD, 20771, USA", "Stellar Astrophysics Centre, Department of Physics &amp; Astronomy, University of Århus, Ny Munkegade 120, 8000, Århus C, Denmark" ]
[ "2015MNRAS.447.2235M", "2015ads..book.....M", "2016AJ....151...49K", "2017A&A...605C...4M", "2017MNRAS.464.2341C", "2017MNRAS.465.1046T", "2018ARA&A..56...83B", "2019ApJ...881..135L", "2019MNRAS.482.1080S", "2020ApJ...889..117L", "2021ApJ...918...27M", "2021MNRAS.502.1703M", "2022MNRAS.514..744M" ]
[ "astronomy" ]
5
[ "stars: horizontal-branch", "stars: atmospheres", "techniques: spectroscopic", "globular clusters: individual: NGC 288", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1927sbss.book.....T", "1979AJ.....84.1511T", "1983ApJ...267..256M", "1985MNRAS.217..305M", "1986PASP...98..609H", "1988ApJ...332..236C", "1995A&A...294...65M", "1995ApJ...453..214P", "1996AJ....112.1487H", "1997A&A...319..109M", "1997A&A...322..256N", "1998ApJ...500..311F", "1999ApJ...517L.135B", "1999ApJ...524..242G", "2000A&A...360..120M", "2000ApJ...528..849B", "2000ApJ...531L..37B", "2000ApJ...533..215C", "2000ApJ...535L..43H", "2000SPIE.4008..129P", "2001PASP..113.1162M", "2003A&A...405..135M", "2003ApJS..149...67B", "2004ASPC..318..402N", "2005A&A...434..235F", "2006A&A...452..493P", "2007A&A...466..667K", "2007A&A...474..505M", "2007AN....328..879S", "2008A&A...477..641K", "2008ApJ...675.1223M", "2008ApJ...684.1159Y", "2009A&A...495..937L", "2009A&A...498..737M", "2009A&A...500.1163Q", "2009A&A...505..139C", "2009A&A...508..695C", "2010A&A...517A..81G", "2010MNRAS.406.2732L", "2010MNRAS.407.1767K", "2010MNRAS.409.1606L", "2011A&A...526A.136M", "2011A&A...529A..60M", "2011A&A...535A.121C", "2011ApJ...733L..45R", "2011ApJ...738L..10M", "2012A&A...547A.109M", "2013A&A...559A.101S", "2013ApJ...775...15P", "2013MNRAS.431.2126M" ]
[ "10.1051/0004-6361/201322953", "10.48550/arXiv.1403.7397" ]
1403
1403.7397_arXiv.txt
\label{sec:intro} Low-mass stars that burn helium in a core of about 0.5\,$M_\sun$ and hydrogen in a shell populate a roughly horizontal region in the optical colour-magnitude diagrams of globular clusters, which has earned them the name horizontal branch ({\bf HB}) stars \citep{ten27}. The hot (or blue) HB stars near an effective temperature of 11\,500\,K are of special interest because they exhibit a number of intriguing phenomena associated with the onset of diffusion. A large photometric survey of many globular clusters by \citet{grca99} demonstrated that the Str\"omgren $u$-brightness of blue HB stars suddenly increases near 11\,500\,K. This $u$-jump is attributed to a sudden increase in the atmospheric metallicity of the blue HB stars to super-solar values that is caused by the {\em radiative levitation of heavy elements}. A similar effect can be seen in broad-band $U, U-V$ photometric data \citep[G1]{fpfd98}. \citet{beco99,beco00} and \citet{mosw00} confirmed with direct spectroscopic evidence that the atmospheric metallicity does indeed increase to solar or super-solar values for HB stars hotter than the $u$-jump. A list of earlier observations of this effect can be found in \citet{moeh01}. Later studies include \citet[M\,3, M\,13, M\,15, M\,68, M\,92, and NGC\,288]{behr03}, \citet[NGC\,1904]{fare05}, and \citet[NGC\,2808]{pare06}. These findings also helped to understand the cause of the low-gravity problem: \citet{crro88} and \citet{mohe95,mohe97} found that hot horizontal branch stars -- when analysed with model spectra of the same metallicity as their parent globular cluster -- show significantly lower surface gravities than expected from evolutionary tracks. Analysing them instead with more appropriate metal-rich model spectra reduces the discrepancies considerably \citep{mosw00}. The more realistic stratified model atmospheres of \citet{hulh00} and \citet{lmhh09} reduce the discrepancies in surface gravity even more (see below for more details). Along with the enhancement of heavy metals, a decrease in the helium abundance by mass $Y$ is observed for stars hotter than $\approx$11\,500\,K, while cooler stars have normal helium abundances within the observational errors. The helium abundance for these hotter stars is typically between 1 and 2 dex smaller than the solar value (e.g. \citealt{behr03}). A trend of the helium abundance relative to $T_{\rm eff}$ was discussed by \citet{momo09,movi12}. Finally, blue HB stars near $\approx$11\,500\,K show a sudden drop in their rotation rates \citep{perc95,bedj00,behr03}, and in some globular clusters (e.g., M\,13) a gap in their HB distribution. The possibility of {\em radiative levitation of heavy elements} and {\em gravitational settling of helium} in HB stars had been predicted long ago by \citet{miva83}, but the discovery of its very sudden onset near 11\,500\,K was a complete surprise. \citet{qcmr09} have shown that helium settling in HB stars cooler than $\approx$11\,500\,K is hampered by meridional flow. In stars hotter than this threshold, helium can settle and the superficial convection zone disappears, so that diffusion might occur in superficial regions of these stars. Recent evolutionary models of HB stars that include atomic diffusion calculated by \citet{miri08,miri11} can reasonably well reproduce the abundance anomalies of several elements observed in blue HB stars. However, these models do not treat the atmospheric region in detail. Instead, they assume an outer superficial mixed zone of approximately 10$^{-7}$M$_\sun$ below which separation occurs. The detection of vertical stratification of some elements, especially iron, in the atmospheres of blue HB stars lends additional support to the scenario of atomic diffusion being active in these regions \citep{klbw07,klbw08,khlb10}, but is at variance with the mixed zone introduced by \citet{miri08,miri11}. \citet{hulh00} and \citet{lmhh09} presented stellar atmosphere models of blue HB stars that include the effect of vertical abundance stratification on the atmospheric structure. These models estimate the vertical stratification of the elements caused by diffusion by assuming that an equilibrium solution (i.e. giving a nil diffusion-velocity) for each species is reached. Assuming a sudden onset of atomic diffusion near 11\,500\,K, these models predict photometric jumps and gaps (\citealt{grca99}, G1 in \citealt{fpfd98}) consistent with observations (see \citealt{lehk10} for more details). The photometric changes, with respect to chemically homogeneous atmosphere models, are due to the modification of the atmospheric structure caused by the abundance stratification. \begin{figure*}[ht] \includegraphics[height=\textwidth,angle=270]{CMD_ngc288_uuy_yby.ps} \caption[]{Blue HB stars in NGC\,288 as observed by \citet{grca99}. The subsample, for which GIRAFFE spectra were observed, is marked by filled circles. The dot-dashed line in the left part marks the $u-y$ colour of the $u$-jump. Stars on the red side of the line do not show evidence for radiative levitation, while stars on the blue side do. The triangles mark overluminous stars around the $u$-jump (from red to blue: 122, 103, 127, 101, 146, 142, 100, 183).}\label{fig:cmd_hb} \end{figure*} The $u$-jump described above can be clearly seen in the colour-magnitude diagram of NGC\,288 shown in Fig.\,\ref{fig:cmd_hb}, where one also finds a group of stars with large (and unexplained) scatter in their $u$ magnitudes and ($u-y$) colours around the jump region (triangles in Fig.\,\ref{fig:cmd_hb}). With maximum errors of 0\fm008 in $u$ and 0\fm003 in $y$ it is unlikely that their positions are caused by photometric errors. The evolutionary status of these stars is unclear. While their bright $u$ magnitudes are suggestive of radiative levitation, the effect would have to be extreme and some of them appear to lie on the cool side of the $u$-jump. Similar groups of unexpectedly bright stars can be found in other globular clusters (see M\,2 and M\,92 in \citealt{grca99}), with NGC\,288 presenting the most pronounced case. We discuss these stars in Sect.\,\ref{sec:lum} in more detail. A colour spread along the red giant branch in NGC\,288 first reported by \citet{yogr08} was identified as a split by \citet{role11}, which was confirmed by \citet{cabr11} and \citet{momi13}. In their excellent review on second and third parameters to explain the horizontal branch morphology, \citet{grca10} estimated that a range in helium abundance of $\Delta Y$~=~0.012 would explain the temperature range of the horizontal branch in NGC\,288. This is consistent with the helium range found more recently from the analysis of main-sequence photometry by \citet{pimi13}. A variation in helium this small is unfortunately too small to be detected by our analysis. \onltab{ \begin{table*}[!h] \caption[]{Target coordinates and photometric data.\label{tab:targ_phot}} \begin{tabular}{rlllllll} \hline \hline ID & $\alpha_{2000}$ & $\delta_{2000}$ & $y$ & $b-y$ & c$_1$ & m$_1$\\ \hline 52 & 00:52:50.77 & $-$26:38:02.9 & 15.357$\pm$0.001 & $+$0.030$\pm$0.003 & $+$1.301$\pm$0.005 & $+$0.131$\pm$0.005\\ 55 & 00:52:47.43 & $-$26:33:14.2 & 15.318$\pm$0.002 & $+$0.094$\pm$0.002 & $+$1.253$\pm$0.004 & $+$0.122$\pm$0.004\\ 60 & 00:52:38.59 & $-$26:37:05.1 & 15.328$\pm$0.003 & $+$0.110$\pm$0.004 & $+$1.202$\pm$0.008 & $+$0.118$\pm$0.007\\ 61 & 00:52:42.80 & $-$26:36:39.2 & 15.312$\pm$0.001 & $+$0.127$\pm$0.002 & $+$1.151$\pm$0.004 & $+$0.121$\pm$0.003\\ 63 & 00:52:47.08 & $-$26:35:25.0 & 15.345$\pm$0.002 & $+$0.094$\pm$0.003 & $+$1.264$\pm$0.005 & $+$0.119$\pm$0.005\\ 70 & 00:52:45.40 & $-$26:35:21.5 & 15.438$\pm$0.002 & $+$0.055$\pm$0.003 & $+$1.288$\pm$0.004 & $+$0.133$\pm$0.004\\ 72 & 00:52:53.92 & $-$26:38:45.6 & 15.421$\pm$0.002 & $+$0.088$\pm$0.004 & $+$1.234$\pm$0.008 & $+$0.125$\pm$0.008\\ 74 & 00:52:35.47 & $-$26:34:24.8 & 15.453$\pm$0.001 & $+$0.055$\pm$0.002 & $+$1.283$\pm$0.005 & $+$0.133$\pm$0.003\\ 79 & 00:52:42.68 & $-$26:34:50.2 & 15.507$\pm$0.002 & $+$0.050$\pm$0.003 & $+$1.307$\pm$0.004 & $+$0.135$\pm$0.005\\ 81 & 00:52:40.77 & $-$26:33:47.7 & 15.569$\pm$0.001 & $+$0.034$\pm$0.002 & $+$1.288$\pm$0.003 & $+$0.137$\pm$0.003\\ 83 & 00:52:32.07 & $-$26:35:46.6 & 15.560$\pm$0.002 & $+$0.022$\pm$0.002 & $+$1.266$\pm$0.007 & $+$0.139$\pm$0.004\\ 86 & 00:52:52.51 & $-$26:34:29.7 & 15.582$\pm$0.001 & $+$0.038$\pm$0.002 & $+$1.284$\pm$0.006 & $+$0.136$\pm$0.004\\ 88 & 00:52:52.10 & $-$26:34:12.1 & 15.639$\pm$0.002 & $+$0.021$\pm$0.003 & $+$1.260$\pm$0.007 & $+$0.131$\pm$0.005\\ 89 & 00:52:37.88 & $-$26:36:35.4 & 15.518$\pm$0.002 & $+$0.042$\pm$0.002 & $+$1.242$\pm$0.006 & $+$0.134$\pm$0.003\\ 90 & 00:52:46.64 & $-$26:39:03.1 & 15.634$\pm$0.002 & $+$0.020$\pm$0.004 & $+$1.253$\pm$0.011 & $+$0.131$\pm$0.009\\ 96 & 00:52:37.64 & $-$26:31:13.5 & 15.735$\pm$0.004 & $+$0.005$\pm$0.004 & $+$1.238$\pm$0.010 & $+$0.134$\pm$0.007\\ 99 & 00:52:29.92 & $-$26:36:07.4 & 15.799$\pm$0.002 & $-$0.004$\pm$0.003 & $+$1.182$\pm$0.009 & $+$0.133$\pm$0.007\\ 100 & 00:52:38.68 & $-$26:35:58.6 & 15.959$\pm$0.002 & $-$0.034$\pm$0.002 & $+$0.858$\pm$0.004 & $+$0.112$\pm$0.004\\ 101 & 00:53:04.16 & $-$26:38:29.8 & 15.911$\pm$0.005 & $-$0.024$\pm$0.009 & $+$0.938$\pm$0.011 & $+$0.116$\pm$0.015\\ 102 & 00:52:51.47 & $-$26:36:26.0 & 15.815$\pm$0.001 & $+$0.007$\pm$0.002 & $+$1.205$\pm$0.004 & $+$0.128$\pm$0.003\\ 103 & 00:52:35.45 & $-$26:39:11.0 & 15.840$\pm$0.003 & $-$0.011$\pm$0.004 & $+$1.063$\pm$0.010 & $+$0.115$\pm$0.007\\ 106 & 00:53:02.55 & $-$26:35:32.9 & 15.822$\pm$0.002 & $+$0.005$\pm$0.003 & $+$1.176$\pm$0.008 & $+$0.127$\pm$0.005\\ 107 & 00:52:32.39 & $-$26:36:30.2 & 15.810$\pm$0.002 & $+$0.008$\pm$0.002 & $+$1.173$\pm$0.009 & $+$0.126$\pm$0.004\\ 111 & 00:53:05.31 & $-$26:32:45.2 & 15.982$\pm$0.003 & $-$0.017$\pm$0.004 & $+$1.084$\pm$0.005 & $+$0.136$\pm$0.005\\ 113 & 00:52:39.48 & $-$26:36:45.7 & 15.866$\pm$0.001 & $+$0.001$\pm$0.002 & $+$1.161$\pm$0.006 & $+$0.125$\pm$0.004\\ 114 & 00:52:38.68 & $-$26:37:49.6 & 15.905$\pm$0.004 & $-$0.002$\pm$0.006 & $+$1.132$\pm$0.015 & $+$0.126$\pm$0.011\\ 115 & 00:52:45.24 & $-$26:37:55.1 & 15.881$\pm$0.001 & $+$0.000$\pm$0.002 & $+$1.170$\pm$0.006 & $+$0.125$\pm$0.005\\ 118 & 00:52:55.50 & $-$26:35:08.2 & 15.902$\pm$0.001 & $-$0.005$\pm$0.002 & $+$1.158$\pm$0.008 & $+$0.124$\pm$0.003\\ 119 & 00:52:56.84 & $-$26:33:44.8 & 15.883$\pm$0.001 & $+$0.002$\pm$0.002 & $+$1.159$\pm$0.007 & $+$0.123$\pm$0.004\\ 120 & 00:53:01.85 & $-$26:37:53.2 & 15.913$\pm$0.004 & $-$0.008$\pm$0.006 & $+$1.136$\pm$0.007 & $+$0.128$\pm$0.010\\ 122 & 00:52:37.78 & $-$26:39:31.6 & 15.880$\pm$0.002 & $-$0.012$\pm$0.004 & $+$1.094$\pm$0.011 & $+$0.127$\pm$0.009\\ 127 & 00:52:39.32 & $-$26:34:31.6 & 15.953$\pm$0.002 & $-$0.014$\pm$0.002 & $+$1.047$\pm$0.004 & $+$0.111$\pm$0.004\\ 142 & 00:52:50.55 & $-$26:36:49.8 & 16.061$\pm$0.002 & $-$0.032$\pm$0.002 & $+$0.876$\pm$0.005 & $+$0.111$\pm$0.004\\ 143 & 00:52:52.77 & $-$26:34:53.0 & 16.003$\pm$0.002 & $-$0.005$\pm$0.003 & $+$1.084$\pm$0.006 & $+$0.119$\pm$0.006\\ 145 & 00:52:40.58 & $-$26:32:48.6 & 15.981$\pm$0.002 & $-$0.013$\pm$0.002 & $+$1.108$\pm$0.004 & $+$0.124$\pm$0.003\\ 146 & 00:52:43.35 & $-$26:37:56.4 & 16.098$\pm$0.002 & $-$0.019$\pm$0.004 & $+$0.878$\pm$0.005 & $+$0.107$\pm$0.006\\ 147 & 00:52:37.26 & $-$26:36:46.9 & 16.033$\pm$0.002 & $-$0.013$\pm$0.003 & $+$1.054$\pm$0.007 & $+$0.127$\pm$0.004\\ 149 & 00:52:33.14 & $-$26:33:44.6 & 16.084$\pm$0.003 & $-$0.020$\pm$0.004 & $+$1.044$\pm$0.009 & $+$0.122$\pm$0.005\\ 151 & 00:52:48.32 & $-$26:32:57.6 & 16.032$\pm$0.002 & $-$0.016$\pm$0.003 & $+$1.074$\pm$0.004 & $+$0.131$\pm$0.004\\ 154 & 00:52:54.89 & $-$26:37:11.7 & 16.058$\pm$0.001 & $-$0.007$\pm$0.002 & $+$1.054$\pm$0.007 & $+$0.112$\pm$0.004\\ 156 & 00:52:50.53 & $-$26:35:12.0 & 16.039$\pm$0.003 & $-$0.004$\pm$0.005 & $+$1.060$\pm$0.007 & $+$0.112$\pm$0.009\\ 157 & 00:52:53.76 & $-$26:39:08.7 & 15.992$\pm$0.004 & $-$0.013$\pm$0.007 & $+$1.073$\pm$0.011 & $+$0.123$\pm$0.013\\ 167 & 00:52:46.42 & $-$26:34:07.7 & 16.093$\pm$0.002 & $-$0.009$\pm$0.003 & $+$1.039$\pm$0.006 & $+$0.112$\pm$0.006\\ 169 & 00:52:46.50 & $-$26:31:30.7 & 16.153$\pm$0.002 & $-$0.023$\pm$0.003 & $+$0.957$\pm$0.006 & $+$0.119$\pm$0.005\\ 176 & 00:52:48.70 & $-$26:34:00.7 & 16.140$\pm$0.002 & $-$0.012$\pm$0.003 & $+$0.984$\pm$0.006 & $+$0.113$\pm$0.006\\ 179 & 00:52:48.17 & $-$26:35:19.9 & 16.236$\pm$0.003 & $-$0.032$\pm$0.006 & $+$0.844$\pm$0.009 & $+$0.108$\pm$0.011\\ 180 & 00:52:50.91 & $-$26:36:09.4 & 16.153$\pm$0.002 & $-$0.014$\pm$0.003 & $+$0.975$\pm$0.004 & $+$0.111$\pm$0.004\\ 183 & 00:52:39.91 & $-$26:37:23.8 & 16.254$\pm$0.002 & $-$0.025$\pm$0.004 & $+$0.784$\pm$0.011 & $+$0.105$\pm$0.009\\ 187 & 00:52:44.71 & $-$26:35:31.4 & 16.307$\pm$0.003 & $-$0.033$\pm$0.005 & $+$0.843$\pm$0.008 & $+$0.116$\pm$0.009\\ 195 & 00:52:27.41 & $-$26:35:58.8 & 16.422$\pm$0.004 & $-$0.050$\pm$0.005 & $+$0.685$\pm$0.008 & $+$0.122$\pm$0.008\\ 196 & 00:52:44.29 & $-$26:35:53.2 & 16.357$\pm$0.002 & $-$0.031$\pm$0.002 & $+$0.760$\pm$0.004 & $+$0.105$\pm$0.004\\ 199 & 00:52:55.57 & $-$26:32:58.7 & 16.425$\pm$0.002 & $-$0.042$\pm$0.002 & $+$0.726$\pm$0.005 & $+$0.113$\pm$0.004\\ 212 & 00:52:59.33 & $-$26:39:00.4 & 16.605$\pm$0.005 & $-$0.068$\pm$0.009 & $+$0.577$\pm$0.013 & $+$0.135$\pm$0.017\\ 213 & 00:52:42.83 & $-$26:31:06.8 & 16.627$\pm$0.005 & $-$0.062$\pm$0.005 & $+$0.564$\pm$0.008 & $+$0.128$\pm$0.007\\ \hline \end{tabular} \end{table*} \setcounter{table}{0} \begin{table*} \caption[]{Target coordinates and photometric data (continued)} \begin{tabular}{rllllll} \hline \hline ID & $\alpha_{2000}$ & $\delta_{2000}$ & $y$ & $b-y$ & c$_1$ & m$_1$\\ \hline 216 & 00:52:54.57 & $-$26:33:20.4 & 16.572$\pm$0.002 & $-$0.051$\pm$0.005 & $+$0.627$\pm$0.011 & $+$0.117$\pm$0.010\\ 221 & 00:52:52.31 & $-$26:35:13.7 & 16.638$\pm$0.002 & $-$0.037$\pm$0.003 & $+$0.584$\pm$0.007 & $+$0.105$\pm$0.006\\ 228 & 00:52:47.36 & $-$26:37:52.6 & 16.761$\pm$0.002 & $-$0.057$\pm$0.004 & $+$0.495$\pm$0.009 & $+$0.116$\pm$0.008\\ 230 & 00:52:24.33 & $-$26:35:23.5 & 16.759$\pm$0.009 & $-$0.059$\pm$0.014 & $+$0.485$\pm$0.029 & $+$0.114$\pm$0.025\\ 231 & 00:52:47.01 & $-$26:36:23.0 & 16.714$\pm$0.003 & $-$0.057$\pm$0.004 & $+$0.552$\pm$0.008 & $+$0.115$\pm$0.006\\ 240 & 00:52:43.69 & $-$26:35:01.7 & 16.761$\pm$0.003 & $-$0.053$\pm$0.005 & $+$0.534$\pm$0.007 & $+$0.113$\pm$0.008\\ 242 & 00:52:45.02 & $-$26:37:35.4 & 16.828$\pm$0.002 & $-$0.064$\pm$0.005 & $+$0.421$\pm$0.007 & $+$0.121$\pm$0.009\\ 243 & 00:52:44.02 & $-$26:35:42.2 & 16.889$\pm$0.003 & $-$0.043$\pm$0.003 & $+$0.445$\pm$0.006 & $+$0.105$\pm$0.005\\ 275 & 00:52:49.39 & $-$26:35:53.7 & 17.030$\pm$0.006 & $-$0.045$\pm$0.009 & $+$0.391$\pm$0.015 & $+$0.098$\pm$0.016\\ 288 & 00:52:49.30 & $-$26:38:19.1 & 17.115$\pm$0.003 & $-$0.058$\pm$0.004 & $+$0.358$\pm$0.006 & $+$0.104$\pm$0.006\\ 292 & 00:53:00.03 & $-$26:36:32.9 & 17.117$\pm$0.003 & $-$0.068$\pm$0.005 & $+$0.406$\pm$0.008 & $+$0.108$\pm$0.008\\ 300 & 00:52:49.11 & $-$26:35:35.1 & 17.106$\pm$0.003 & $-$0.048$\pm$0.007 & $+$0.362$\pm$0.009 & $+$0.100$\pm$0.013\\ 304 & 00:52:48.60 & $-$26:33:17.1 & 17.173$\pm$0.003 & $-$0.054$\pm$0.005 & $+$0.307$\pm$0.011 & $+$0.107$\pm$0.009\\ 347 & 00:52:50.31 & $-$26:38:30.0 & 17.532$\pm$0.003 & $-$0.066$\pm$0.005 & $+$0.275$\pm$0.009 & $+$0.102$\pm$0.009\\ \hline \end{tabular} \end{table*} }
\begin{itemize} \item The atmospheric parameters and masses of the hot HB stars in NGC\,288 show a behaviour also seen in other clusters for temperatures between 9\,000\,K and 14\,000\,K. Outside this temperature range, however, they follow the results found for such stars in $\omega$\,Cen. \item The abundances derived from our observations are for most elements within the abundance range expected from evolutionary models that include the effects of atomic diffusion and assume a surface mixed mass of 10$^{-7}$M$_\sun$, as determined previously for sdB stars and other clusters. The exceptions are helium, which is more deficient than expected, and phosphorus, which is substantially more abundant than predicted. The abundances predicted by stratified model atmospheres, which were not adjusted to observations, are generally significantly more extreme than observed, except for magnesium. \item When the observed spectra were analysed with stratified model spectra, the HB stars were moved to higher temperatures, surface gravities, and masses. Since the equilibrium abundances led to excessive adjustments, more realistic abundance gradients may well lead to models that locate the HB stars between the TAHB and ZAHB (see Sect.\,\ref{sec:lineprofstrat}). Model atmospheres including such improvements are needed to answer this question. \item Five of the eight overluminous stars around the $u$-jump that we observed do not deviate substantially from the other HB stars in the same temperature range in any of the parameters we determined. We are therefore at a loss to explain their brighter luminosities. The remaining three overluminous stars do show lower surface gravities, as would be expected if they evolved away from the HB. However, they pose a substantial problem for evolutionary time-scales, because one would expect to find approximately100 HB stars close to the ZAHB between about 11\,000\,K and 14\,000\,K (corresponding to a range of roughly 0.5 -- 1.05 in $u-y$ in Fig.\,\ref{fig:cmd_hb}) for each of the evolved stars, which is clearly not the case. \item All overluminous stars show the same abundance behaviour as the majority of the stars in the respective temperature range. This result supports the general assumption that the abundances of HB stars redward of the $u$-jump are not affected by diffusion, irrespective of their evolutionary status. \end{itemize} Evolution models including diffusion and stratified model atmospheres both predict higher-than-observed abundances for many elements affected by radiative levitation. Evolution models can be adjusted to reproduce the observed abundances by introducing an ad hoc defined mixed zone, which is potentially inconsistent with the observed vertical stratification of at least some elements in the atmosphere, however. For stratified model atmospheres it is currently unclear whether a more limited abundance stratification, which would provide a better description of the observed abundances, would still be able to explain the photometric anomalies around the $u$-jump. Our results show definite promise towards solving the long-standing problem of surface gravity and mass discrepancies for hot HB stars, but much work is still needed to arrive at a self-consistent solution.
14
3
1403.7397
Context. NGC 288 is a globular cluster with a well-developed blue horizontal branch (HB) covering the u-jump that indicates the onset of diffusion. It is therefore well suited to study the effects of diffusion in blue HB stars. <BR /> Aims: We compare observed abundances with predictions from stellar evolution models calculated with diffusion and from stratified atmospheric models. We verify the effect of using stratified model spectra to derive atmospheric parameters. In addition, we investigate the nature of the overluminous blue HB stars around the u-jump. <BR /> Methods: We defined a new photometric index sz from uvby measurements that is gravity-sensitive between 8000 K and 12 000 K. Using medium-resolution spectra and Strömgren photometry, we determined atmospheric parameters (T<SUB>eff</SUB>, log g) and abundances for the blue HB stars. We used both homogeneous and stratified model spectra for our spectroscopic analyses. <BR /> Results: The atmospheric parameters and masses of the hot HB stars in NGC 288 show a behaviour seen also in other clusters for temperatures between 9000 K and 14 000 K. Outside this temperature range, however, they instead follow the results found for such stars in ω Cen. The abundances derived from our observations are for most elements (except He and P) within the abundance range expected from evolutionary models that include the effects of atomic diffusion and assume a surface mixed mass of 10<SUP>-7</SUP> M<SUB>⊙</SUB>. The abundances predicted by stratified model atmospheres are generally significantly more extreme than observed, except for Mg. When effective temperatures, surface gravities, and masses are determined with stratified model spectra, the hotter stars agree better with canonical evolutionary predictions. <BR /> Conclusions: Our results show definite promise towards solving the long-standing problem of surface gravity and mass discrepancies for hot HB stars, but much work is still needed to arrive at a self-consistent solution. <P />Based on observations with the ESO Very Large Telescope at Paranal Observatory, Chile (proposal ID 71.D-0131).Tables 1 and 2 are available in electronic form at <A href="http://www.aanda.org/10.1051/0004-6361/201322953/olm">http://www.aanda.org</A>The observed abundances plotted in Fig. 8 are only available at the CDS via anonymous ftp to <A href="http://cdsarc.u-strasbg.fr">http://cdsarc.u-strasbg.fr</A> (ftp://130.79.128.5) or via <A href="http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/565/A100">http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/565/A100</A>
false
[ "blue HB stars", "hot HB stars", "stratified atmospheric models", "stratified model spectra", "stratified model atmospheres", "evolutionary models", "atomic diffusion", "stellar evolution models", "diffusion", "such stars", "HB", "atmospheric parameters", "surface gravities", "canonical evolutionary predictions", "anonymous ftp", "the blue HB stars", "observed abundances", "the hot HB stars", "ω Cen", "other clusters" ]
8.556873
8.85758
140
814129
[ "Ishida, Hiroyuki", "Kusakabe, Motohiko", "Okada, Hiroshi" ]
2014PhRvD..90h3519I
[ "Effects of long-lived 10 MeV-scale sterile neutrinos on primordial elemental abundances and the effective neutrino number" ]
29
[ "Department of Physics, Tohoku University, Sendai 980-8578, Japan", "School of Liberal Arts and Science, Korea Aerospace University, Goyang 412-791, Korea and Department of Physics, Soongsil University, Seoul 156-743, Korea", "School of Physics, KIAS, Seoul 130-722, Korea" ]
[ "2014PhRvD..90d5009K", "2014PhRvD..90h5018C", "2015PhRvD..91j3007P", "2015PhRvD..91j4023K", "2015PhRvL.114i1101P", "2016JCAP...08..022S", "2016NuPhB.908..336G", "2016NuPhB.911..955G", "2016PhRvD..93d3511K", "2016PhRvD..93g3008E", "2017IJMPE..2641004K", "2017JHEP...06..050F", "2017PhRvD..95d3524L", "2018PhRvD..97d3005J", "2018PhRvD..98h3503H", "2019ApJ...872..172L", "2019JHEP...01..074F", "2019PhLB..79834980A", "2019PhRvD..99d3505K", "2019PhRvD.100e3006B", "2019PhRvD.100g3011B", "2019arXiv191211453W", "2020JCAP...12..048K", "2020JHEP...07..179C", "2020PhRvD.101h3010L", "2021JCAP...05..042L", "2021JCAP...07..027C", "2023MNRAS.526.2920V", "2024Univ...10..115O" ]
[ "astronomy", "physics" ]
5
[ "26.35.+c", "95.35.+d", "98.80.Cq", "98.80.Es", "Big Bang nucleosynthesis", "Dark matter", "Particle-theory and field-theory models of the early Universe", "Observational cosmology", "Astrophysics - Cosmology and Nongalactic Astrophysics", "High Energy Physics - Phenomenology" ]
[ "1948PhRv...73..803A", "1963PhRv..131.2578P", "1963PhRv..131.2582T", "1963PhRvL..10..531C", "1968PhRv..167.1159J", "1969JETPL...9..184S", "1973PThPh..49..652K", "1977PhLB...66..202S", "1977PhLB...67..421M", "1978ApJ...221..327D", "1979MNRAS.188P..15L", "1980PThPh..64.1103Y", "1981PhRvC..23..645L", "1982A&A...115..357S", "1982PhRvD..25.1481K", "1983NuPhA.395..119W", "1985NuPhB.259..175E", "1986PhLB..174...45F", "1987ApJ...319..531K", "1988ApJ...330..545D", "1988ApJ...331...19S", "1988ApJ...331...33S", "1988PhRvD..37.3441R", "1990ApJS...73...21D", "1990acr..book.....B", "1990eaun.book.....K", "1992NuPhB.373..399E", "1992STIN...9225163K", "1993ApJS...85..219S", "1993PhRvL..70.2661H", "1995ApJ...452..506K", "1995PThPh..93..879K", "1995PhRvD..51.4134P", "1996ApJ...473..576F", "1996RPPh...59.1493S", "1997PhRvD..56.2970J", "1999ApJ...527..180P", "1999NuPhA.656....3A", "2000ApJ...530L..57R", "2000NuPhB.580..331D", "2000NuPhB.590..562D", "2000PhRvL..84.3248J", "2001PhRvD..63j3502K", "2002ApJ...574..398P", "2002Natur.415...54B", "2003ApJS..148..175S", "2003NuPhA.729..337A", "2003PhRvD..67j3521C", "2003PhRvD..68f3504F", "2004ApJ...615L..33M", "2004PhRvD..70d3506H", "2005ApJ...619..538R", "2005JHEP...11..028K", "2005NIMPA.552..329K", "2005Natur.434..871F", "2005PhLB..605...72S", "2005PhLB..619...30E", "2005PhRvC..72d4004S", "2005PhRvD..71b1302M", "2005PhRvD..71h3502K", "2005PhRvD..72d3522I", "2006ApJ...639..897A", "2006ApJ...644..229A", "2006JCAP...10..016S", "2006PhRvD..74b3526K", "2006PhRvD..74j3509J", "2007A&A...462..851B", "2007A&A...465..587S", "2007ApJ...671..402K", "2007ApJS..170..377S", "2007JHEP...01..091A", "2007JHEP...10..015G", "2008PhRvC..78f4614C", "2009A&A...503..545L", "2009A&A...505L..13G", "2009ApJ...698.1803A", "2009ApJ...707..916F", "2009PhRvD..79l3513K", "2009PhRvL.102w2502D", "2009PhRvL.103o9903D", "2010A&A...519L...3M", "2010A&A...522A..26S", "2010ApJ...710L..67I", "2010ApJS..189..240C", "2010JCAP...05..003A", "2010PhRvC..82c5501S", "2011ARNPS..61...47F", "2011ApJS..192...16L", "2011JHEP...04..011A", "2011MNRAS.412...81M", "2011arXiv1110.6479F", "2012A&A...539A.157M", "2012ApJ...744..158C", "2012ApJ...751L...6A", "2012JHEP...07..112A", "2012MNRAS.419.2195M", "2012MNRAS.425.2477P", "2012MSAIS..22...35A", "2012PhRvC..86c2801C", "2012PhRvD..86a0001B", "2012PhRvD..86f3003K", "2013A&A...554A..96L", "2013ApJS..208...19H", "2013JHEP...03..125A", "2013NuPhA.908....1B", "2013PhLB..718..704K", "2013PhRvD..87h5045K", "2013arXiv1307.6955C", "2014A&A...571A..16P" ]
[ "10.1103/PhysRevD.90.083519", "10.48550/arXiv.1403.5995" ]
1403
1403.5995_arXiv.txt
} Big bang nucleosynthesis (BBN) model \cite{Alpher1948} successfully explains primordial light element abundances inferred from astronomical observations (e.g. \cite{Coc:2011az,Coc:2013eea}) if the cosmological baryon density determined from the power spectrum of cosmic microwave background (CMB) radiation measured with the Wilkinson Microwave Anisotropy Probe (WMAP) \cite{Spergel:2003cb,Spergel:2006hy, Larson:2010gs,Hinshaw:2012fq} or Planck \cite{Ade:2013zuv} is adopted. An apparent discrepancy, however, exists between observational and theoretical $^7$Li abundances. Spectroscopic observations of metal-poor stars (MPSs) indicate a plateau for the abundance ratio, $^7$Li/H$=(1-2) \times 10^{-10}$, with small error bars as a function of metallicity for [Fe/H]$>-3$ \footnote{[A/B]$=\log(n_{\rm A}/n_{\rm B})-\log(n_{\rm A}/n_{\rm B})_\odot$, where $n_i$ is the number density of element ($i$=A and B), and the subscript $\odot$ indicates the solar value.} in the Galaxy~\cite{Spite:1982dd,Ryan:1999vr,Melendez:2004ni,Asplund:2005yt, bon2007,shi2007,Aoki:2009ce,Hernandez:2009gn,Sbordone:2010zi,Monaco:2011sd,Mucciarelli:2011ts,Aoki:2012wb,Aoki2012b} and $\omega$ Centauri accreted by the Galaxy~\cite{Monaco:2010mm} \footnote{Average stellar Li abundances in metal-poor globular clusters (GCs; e.g. \cite{Hernandez:2009gn,Mucciarelli:2010gz}) are larger than those in metal-poor halo stars. The Li abundance in GC M4 turn-off stars has been determined to be log($^7$Li/H)$=-12+(2.30\pm 0.02 \pm 0.10)$ \cite{Mucciarelli:2010gz}, while the Li abundance in halo dwarf stars is log($^7$Li/H)$=-12+(2.199\pm 0.086)$ \cite{Sbordone:2010zi}. However, these abundances are consistent with each other within the uncertainties for the moment (see fig. 3 of Ref. \cite{Mucciarelli:2010gz}). The possible systematic difference in the Li abundances in GCs and Galactic halo should be studied further in future.}. The plateau abundance is $\sim 3-4$ times smaller than that predicted in standard BBN (SBBN) model (e.g., $^7$Li/H=$5.24 \times 10^{-10}$~\cite{Coc:2011az}; see Ref. \cite{Coc:2013eea} for theoretical light element abundances for the baryon density from Planck \cite{Ade:2013zuv}). Recent observations also indicate a break of this plateau shown as small Li abundances with large dispersion at lower metallicities of [Fe/H]$<-3$ \cite{Frebel:2005ig,Aoki:2005gn,bon2007,Aoki:2009ce,Hernandez:2009gn,Sbordone:2010zi,Aoki2012b}. Therefore, it seems that we need a mechanism for a metallicity-independent depletion from the primordial abundance to the plateau abundance and also another for a metallicity-dependent depletion from the plateau abundance. In this paper, we focus on the former universal depletion by cosmological processes. This Li problem (see Ref. \cite{Fields:2011zzb} for a review) shows that some physical processes have reduced the primordial Li abundance in some epoch during or after BBN. Standard stellar model suggests very small depletions of Li isotopes in surfaces of MPSs \cite{Deliyannis1990}. The $^7$Li/H abundances of MPSs observed today are then approximately interstellar abundances when the MPSs formed. Nonstandard processes such as the rotationally induced mixing \cite{Pinsonneault:1998nf,Pinsonneault:2001ub}, and the turbulent mixing~\cite{Richard:2004pj,Korn:2007cx,Lind:2009ta} have been suggested to reduce the $^7$Li abundance in stellar atmospheres. In the former model, a large depletion factor does not realize simultaneously with a small dispersion in stellar Li abundances after the depletion. The depletion factor is then constrained to be small. In the latter model, a depletion of a factor of $1.6-2.0$ \cite{Richard:2004pj} is predicted although it is still unclear if this mechanism can deplete Li abundances of all MPSs rather uniformly. Nonstandard BBN, on the other hand, may be responsible for the Li problem at least partially. We note that $^7$Be is produced more than $^7$Li in SBBN model with the Planck baryon density. The $^7$Be nuclei are then converted to $^7$Li nuclei via recombination with electron followed by the electron capture, i.e., $^7$Be + e$^-\rightarrow ^7$Li +$ \nu_e$. Therefore, the Li problem is alleviated if some exotic processes could destroy $^7$Be. One of solutions to the Li problem is an injection of nonthermal photon with energy of $\sim 2$ MeV which can destroy $^7$Be but not Deuterium (D) as calculated in Ref. \cite{Kusakabe:2013sna}. If a long-lived exotic particle radiatively decays after BBN, nonthermal photons can disintegrate background thermal nuclei, and light element abundances change \cite{Lindley1979MNRAS.188P..15L}. If the energy of the photon emitted at the decay is much larger than ${\cal O}(10)$ MeV, all of light nuclei are disintegrated by nonthermal photons~\cite{Lindley1979MNRAS.188P..15L,Ellis:1984er, Dimopoulos:1987fz,1992NuPhB.373..399E, Kawasaki:1994af,Kawasaki:1994sc,Jedamzik:1999di,Kawasaki:2000qr,Cyburt:2002uv, Kawasaki:2004qu,Ellis:2005ii,Jedamzik:2006xz,Kusakabe:2006hc}. In this case, therefore, the Li problem can not be solved (e.g., \cite{Ellis:2005ii,Kusakabe:2006hc}). Therefore, the energy of emitted photon for the $^7$Be destruction is limited to a narrow range \cite{Kusakabe:2013sna}. A similar $^7$Be destruction would occur if a long-lived sterile neutrino decays into energetic electron and positron. We then study cosmological effects of this decay channel in this paper. From the theoretical point of view of the extended Minimal Standard Model (MSM) of particle physics, right-handed neutrinos introduced as sterile neutrinos provide an elegant mechanism for the generation of tiny active neutrino masses; so-called canonical seesaw mechanism \cite{Minkowski:1977sc,Yanagida1979,Yanagida:1980xy,Gell-Mann1979}. If their masses are so heavy (more than $10^9~{\rm GeV}$), they also explain the origin of the baryon asymmetry of the universe; so-called leptogenesis scenario \cite{Fukugita:1986hr}. Even if the sterile neutrinos have masses below electroweak (EW) scale, however, there exist other phenomenological effects without lacking the success of the seesaw mechanism. Several possibilities have been investigated regarding a detectability of the sterile neutrinos, and the case in which sterile neutrinos are lighter than light mesons (e.g., pion or kaon) has been studied especially in detail \cite{Asaka:2011pb} (for recent study, see \cite{Asaka:2012bb} and references therein). The allowed sterile neutrino masses are smaller than the pion mass $\sim$ 140 MeV, which are not consistent with recent results of neutrino oscillation experiments, unless one assumes that their lifetimes are longer than $\sim 0.1$ s. % In addition, another constraint has been derived from a study of BBN by a comparison between theoretical and observational abundances of $^4$He. The upper limit on the lifetime of $\sim 0.1$ s was derived when a relic abundance of sterile neutrino is fixed as given in Ref. \cite{Dolgov:2000pj,Dolgov:2000jw}. However, this constraint depends on the relic abundance. In this paper, we take into account the possibility that the abundance is smaller than the simple estimate \cite{Dolgov:2000pj,Dolgov:2000jw}. In this case a longer lifetime of the sterile neutrino is allowed. In this paper, we comprehensively investigate cosmological effects of a long-lived sterile neutrino with mass $\sim 10$ MeV. In Sec. \ref{sec2}, we assume a decay of a sterile neutrino in the early universe, and describe our calculation method, and formulations of (1) the spectra of electrons and positrons generated at the decay, (2) those of primary photons induced by the energetic electrons and positrons, and (3) the nonthermal nucleosynthesis triggered by the energetic photons. In Sec. \ref{sec3}, revised cross sections for photodisintegration of $^7$Be and $^7$Li are described. In Sec. \ref{sec4}, effects of the decaying sterile neutrino on the cosmic thermal history, the effective neutrino number, and the baryon-to-photon ratios are formulated. In Sec. \ref{sec5}, observational constraints on primordial light element abundances, the effective neutrino number, and the baryon-to-photon ratio adopted in this paper are summarized. In Sec. \ref{sec6}, we show calculated energy spectra of electrons and positrons emitted at the decay, energy spectra of photons produced via the inverse Compton scattering of background photons by the electron and positron, and photon injection spectra resulting from electromagnetic cascade showers. Time evolutions of light element abundances, the baryon-to-photon ratio, and thermal and nonthermal neutrino energy densities are then consistently calculated with nonthermal photodisintegrations of nuclei taken into account. An impact of revised cross sections of $^7$Be and $^7$Li photodisintegrations is also shown. In Sec. \ref{sec7}, we discuss the relic abundance of the sterile neutrino before its decay. We also comment on a possible dilution of the sterile neutrino in the early universe, effects of the sterile neutrino mixing to active neutrinos of different flavors, and experimental constraints from the pion decay and the supernova luminosity. In Sec. \ref{sec8}, this study is summarized. In Appendix \ref{app1}, extensive formulae of the sterile neutrino decay are derived. We adopt natural units of $\hbar=c=k_{\rm B}=1$, where $\hbar$ is the reduced Planck constant, $c$ is the speed of light, and $k_{\rm B}$ is the Boltzmann constant. We also adopt notation of $A$($a$, $b$)$B$ for a reaction $A+a \rightarrow b + B$.
\label{sec7} \subsection{relic abundance of sterile neutrino}\label{sec7a} It is assumed that the sterile neutrino has a mass $M_\nuh$ after the EW phase transition. Depending on the mixing angle, the sterile neutrino react with standard model particles mainly via the following weak reactions: \begin{eqnarray} \nuh + \nu_e &\rightarrow& f +\bar{f} \label{process_1}\\ \nuh + f &\rightarrow& \nu_e +f \\ \nuh + \bar{f} &\rightarrow& \nu_e +\bar{f} \\ \nuh + e^+ &\rightarrow& \bar{e_n}^+ + \nu_{e_n} \\ \nuh + e^+ &\rightarrow& u_n^{+2/3} + \bar{d_{n^\prime}}^{+1/3} \\ \nuh + e_n^- &\rightarrow& e^- + \nu_{e_n} \\ \nuh + \bar{\nu}_{e_n} &\rightarrow& e^- + \bar{e_n}^+ \\ \nuh + \bar{u_n}^{-2/3} &\rightarrow& e^- + \bar{d_{n^\prime}}^{+1/3} \\ \nuh + d_n^{-1/3} &\rightarrow& e^- + u_{n^\prime}^{+2/3}, \label{process_9} \end{eqnarray} where $f$ is any fermion, i.e., charged leptons $e_n$ [$e^-$ ($n=1$), $\mu^-$ ($n=2$), and $\tau^-$ ($n=3$)], neutrinos $\nu_{e_n}$ and up-type quarks $u_n$ [$u$ ($n=1$), $c$ ($n=2$), and $t$ ($n=3$)], and down-type quarks $d_n$ [$d$ ($n=1$), $s$ ($n=2$), and $b$ ($n=3$)]. In the charged current reactions, probabilities of producing respective flavors ($n_n \leftrightarrow d_{n^\prime}$) are described by the Cabibbo-Kobayashi-Maskawa matrix \cite{Cabibbo:1963yz,Kobayashi:1973fv}. When the weak interaction rate becomes smaller than the Hubble expansion rate, the abundance of the sterile neutrino freezes out from equilibrium. Thereafter, the ratio between the $\nuh$ number density and the entropy density $Y_\nuh\equiv n_\nuh/s$ does not change \footnote{We note that this ratio $Y_\nuh$ is measured in a unit different from that of nuclear mole fractions $Y_A$ introduced in Sec. \ref{sec2g}.}. Weak interaction rates of sterile neutrinos, $\nuh$'s, with weakly interacting standard model particles after the EW phase transition scale \cite{Dolgov:2000pj} as \begin{equation} \Gamma \sim G_{\rm F}^2 \Theta^2 T^5, \end{equation} where $G_{\rm F}$ is the Fermi constant, $\Theta\ll 1$ is the mixing angle. The ratio between the rates and cosmic expansion rate, $H$, is then given \cite{Fuller:2011qy} by \begin{eqnarray} \frac{\Gamma}{H} &\sim& G_{\rm F}^2 \Theta^2 T^5 \left( \frac{2\pi^{3/2}}{3 \sqrt[]{\mathstrut 5}} \frac{g_\ast^{1/2} T^2}{M_{\rm Pl}} \right)^{-1} \nonumber\\ &=& \frac{3 \sqrt[]{\mathstrut 5}}{2\pi^{3/2}} \frac{M_{\rm Pl}G_{\rm F}^2 \Theta^2 T^3}{g_\ast^{1/2}} \nonumber\\ &=& 9.69 \times 10^{7} \left(\frac{\Theta}{10^{-3}} \right)^2 \left(\frac{g_\ast}{106.75}\right)^{-1/2} \left(\frac{T}{100~{\rm GeV}}\right)^3 \nonumber\\ &=& 1.00 \left(\frac{\Theta}{10^{-3}} \right)^2 \left(\frac{g_\ast} {63.75}\right)^{-1/2} \left(\frac{T}{0.2~{\rm GeV}}\right)^3. \label{eq_temp_freezeout} \end{eqnarray} In the last line of this equation, the statistical DOF of the sterile neutrino, i.e., $\Delta g_\ast=2\times 7/8$ were added to the value of $g_\ast=61.75$ at $T=200$ MeV in the standard model. A sterile neutrino with mixing angle $\Theta\sim 10^{-3}$ would thus freeze out from equilibrium at temperature $T\sim 200$ MeV. The relic abundance of $\nuh$ is, therefore, given by the abundance fixed at $T\sim 200$ MeV. The lifetime of the sterile neutrinos is roughly given [cf. Eqs. (\ref{eq_a9}) and (\ref{eq_a21})] by \begin{eqnarray} \Gamma(\nuh-{\rm decay}) &\sim& \frac{G_{\rm F}^2 \Theta^2 M_\nuh^5}{192 \pi^3} \nonumber\\ &=& 1.87 \times 10^{-5}~{\rm s}^{-1} \left( \frac{\Theta}{10^{-3}}\right)^2 \left( \frac{M_\nuh} {14~{\rm MeV}}\right)^5. \label{gamma_approx} \end{eqnarray} Therefore, if a sterile neutrino with a mass $M_\nuh \gtrsim 14$ MeV decays $\sim 10^4 -10^5$ s after the big bang and reduces the primordial $^7$Li abundance, the mixing angle would be $\Theta \sim 10^{-3}$. The time evolution of the $\nuh$ abundance has been calculated \cite{Kolb:1981cx,Scherrer:1987rr,Dolgov:2000pj}. In Refs. \cite{Kolb:1981cx,Scherrer:1987rr}, however, the maximal mixing angle $\Theta ={\mathcal O}(1)$ is implicitly assumed, and the dependence on the mixing angle $\Theta$ is not considered. In addition, the authors took into account only the annihilation $\nuh +\bar{\nuh}$ \cite{Dicus:1977qy}, which is negligibly weaker than the reactions Eqs. (\ref{process_1})-(\ref{process_9}) when $\Theta \ll 1$. In Ref. \cite{Dolgov:2000pj}, on the other hand, a dedicated calculation has been performed. However, the authors only focused on shorter $\nuh$ lifetimes of $\tau_\nuh={\mathcal O}(0.1)$ s, which correspond to relatively large values of the mixing angle $\Theta >{\mathcal O}(10^{-3})$ compared with those considered in this paper. Depending on the mixing angle, the weak reaction freeze-out of the sterile neutrino occurs in various epochs with different values of $g_\ast$. Perhaps the sterile neutrino never experiences the weak reaction equilibrium after the EW phase transition. In general, the $\nuh$ relic abundance can sensitively depend on the evolution of sterile neutrino mass during the EW phase transition, which differs in different models of the sterile neutrino. Precise calculations of the $\nuh$ relic abundance should be performed in detail. However, they are beyond the scope of this paper. We estimate the freeze-out abundance of the sterile neutrino as a function of $M_\nuh$ and $\tau_\nuh$ as follows. For a given set of ($M_\nuh$, $\tau_\nuh$), a corresponding $\Theta$ value is derived with Eq. (\ref{gamma_approx}). The temperature satisfying Eq. (\ref{eq_temp_freezeout}) is then derived with the $\Theta$ value. This temperature is defined as the freeze-out temperature $T_{\rm F}$. An approximate value of the freeze-out abundance of $\nuh$ is given by the equilibrium abundance at $T_{\rm F}$. For the $T_{\rm F}(M_\nuh, \tau_\nuh)$ value, the freeze-out abundance is given by the equilibrium abundance $Y_{\nuh,{\rm EQ}}(M_\nuh, T_{\rm F})$ using the following equations. The equilibrium number density of a fermion is given \cite{kolb1990} by \begin{equation} n_{i,{\rm EQ}}(m_i,T)= \frac{g_iT^3}{2\pi^2} h(m_i/T), \label{n_eq} \end{equation} where $m_i$ and $g_i$ are the mass and statistical DOF, respectively, of the fermion $i$, and $h(x)$ is a function given by \begin{equation} h(x)= \int_x^\infty \frac{\left( \epsilon^2 -x^2 \right)^{1/2} \epsilon}{\exp\left( \epsilon \right) +1} d\epsilon. \label{h_plus} \end{equation} In the nonrelativistic limit, the function $h$ has the limit value of $h \rightarrow 3\zeta(3)/2$. The entropy density is given \cite{kolb1990} by \begin{equation} s(T)= \frac{2 \pi^2}{45} g_{\ast{\rm S}} T^3. \label{eq_entropy_density} \end{equation} From Eqs. (\ref{n_eq}) and (\ref{h_plus}), the abundance ratio is given by \begin{equation} Y_{i,{\rm EQ}}(m_i,T) \equiv \frac{n_{i,{\rm EQ}}(m_i,T)}{s(T)}= \frac{45 g_i}{4\pi^4 g_{\ast{\rm S}}} h(m_i/T). \label{eq_entropy_density} \end{equation} Figure \ref{fig_relic} shows massless DOFs in terms of energy and entropy, i.e., $g_\ast$ and $g_{\ast{\rm S}}$, respectively, as a function of photon temperature $T$. Solid lines for massless DOFs correspond to the standard model plus a sterile neutrino of mass $M_\nuh=14$ MeV and statistical DOF of $g_\nuh=2$, while dashed lines correspond to the standard model. Also shown is the equilibrium abundance ratio of a sterile neutrino $Y_{\nuh,{\rm EQ}}$ in the model with the sterile neutrino. The massless DOFs are calculated as in Ref. \cite{kolb1990} based on the latest data on particle mass \cite{Beringer:1900zz}. It is assumed that the quark hadron transition occurs suddenly at temperature $T_{\rm C}=150$ MeV. Above the temperature, quarks are taken into account in the DOFs. Below the temperature, on the other hand, contributions of only hadrons are included and those of quarks are neglected. We only take into account DOFs of charged and neutral pions at $T<T_{\rm C}$ since they are only relativistic hadrons. \begin{figure} \begin{center} \includegraphics[width=8.0cm,clip]{fig_relic.eps} \caption{Massless DOFs in terms of energy and entropy, i.e., $g_\ast$ and $g_{\ast{\rm S}}$, respectively, as a function of photon temperature $T$. Solid lines for massless DOFs correspond to the standard model plus a sterile neutrino of mass $M_\nuh=14$ MeV and statistical DOF of $g_\nuh=2$, while dashed lines correspond to the standard model. Also shown is the equilibrium abundance ratio of a sterile neutrino $Y_{\nuh,{\rm EQ}}$ in the model with the sterile neutrino. \label{fig_relic}} \end{center} \end{figure} The massless DOFs in the model with $\nuh$ is larger than those in the model without $\nuh$ by about two because of the statistical DOF of sterile neutrino. As the temperature decreases, weak bosons, heavy quarks and leptons become nonrelativistic, and the DOFs become small. At the quark hadron transition temperature $T=T_{\rm C}$, DOFs of quarks and gluons disappear, and the DOFs drastically decreases. The equilibrium abundance $Y_{\nuh,{\rm EQ}}$ increases as the temperature increases since it is proportional to the $g_{\ast{\rm S}}(T)$ value. At $T_{\rm C}$, the abundance significantly increases. At the lowest temperature of $T\lesssim 20$ MeV, the sterile neutrino start to be nonrelativistic. The equilibrium abundance then decreases from this temperature. If a light sterile neutrino with the mass $M_\nuh ={\cal O}(10)$ MeV survives during BBN epoch, its number density must have diluted between its weak freeze-out [$T\sim {\cal O}(100)$ MeV] and the BBN epoch [$T\sim {\cal O}(0.1)$ MeV] in order to avoid a large change of the baryon-to-photon ratio associated with the $\nuh$ decay (see Sec. \ref{sec7b}). For example, we consider the case of $M_\nuh=14$ MeV and $\zeta_{\nuh\rightarrow e}=3\times 10^{-7}$ GeV. This assumption corresponds to the energy ratio $\zeta_{\nuh\rightarrow e}/\zeta_{\nuh\rightarrow \nu}=0.313$ [Eqs. (\ref{eq_ratio_nue}) and (\ref{eq_zeta_ratio})], and the total energy injection of $\zeta_{\nuh}=\zeta_{\nuh\rightarrow e}+\zeta_{\nuh\rightarrow \nu}=1.26\times 10^{-6}$ GeV. This energy injection is realized by the decay of sterile neutrino with the mass $M_\nuh=14$ MeV and the number ratio $Y_\nuh=1.28\times 10^{-5}$, where we used a relation, \begin{equation} \zeta_{\nuh} =\frac{n_\nuh}{s} \frac{s}{n_\gamma} M_\nuh =7.04Y_\nuh M_\nuh, \label{relation_zeta_Y} \end{equation} where the ratio $s/n_\gamma=7.04$ should be measured after the cosmological $e^\pm$ annihilation. However, the freeze-out abundance is $Y_\nuh=6.56\times 10^{-3} (g_{\ast{\rm S}}/63.5)^{-1}$. Therefore, the sterile neutrino needs to be diluted by a factor of several hundreds. \subsection{dilution of sterile neutrino}\label{sec7b} It is shown that the decays of heavier sterile neutrinos into standard model particles can realize a dilution of the light sterile neutrino ($M_\nuh\sim 14$ MeV) although our naive estimate~\cite{Asaka:2006nq} suggests that a decay of a heavier sterile neutrino would result in a dilution factor smaller than required for the appropriate abundance by some factor at least. We assume that one of heavier sterile neutrinos, i.e., $\nuh_2$, predominantly contributes to the dilution or entropy production. In addition, it is assumed that the heavy neutrino dominates in terms of energy density in its decay epoch. Supposing that $\nuh_2$ decays into relativistic leptons and quarks which are thermalized rapidly with respect to the cosmic expansion time scale, the energy density of relativistic species after the decay is $\rho_{\rm R} = \frac{\pi^2}{30} g_\ast T_{\rm RH}^4$, where $T_{\rm RH}$ is the reheating temperature. This energy density is equal to the energy density of $\nuh_2$ before the decay. The ratio of the entropy per comoving volume at the epoch long after the decay (aft) to that long before the decay (bef) is given \cite{kolb1990} by \begin{eqnarray} \frac{S_{\rm aft}}{S_{\rm bef}} &=& \frac{g_{\ast{\rm S}}(T_{\rm aft}) a_{\rm aft}^3 T_{\rm aft}^3}{g_{\ast{\rm S}}(T_{\rm bef}) a_{\rm bef}^3 T_{\rm bef}^3} \nonumber\\ &\simeq & 1.83\langle g_\ast^{1/3} \rangle^{3/4} \frac{m_{\nuh_2} Y_{\nuh_2} \tau_{\nuh_2}^{1/2}}{M_{\rm Pl}^{1/2}} \nonumber\\ &= & 8.25 \times 10^1 \left( \frac{\langle g_\ast^{1/3} \rangle^{3/4}}{104^{1/4}} \right) \left( \frac{m_{\nuh_2}}{100~{\rm GeV}} \right) \left( \frac{Y_{\nuh_2}}{4.00\times 10^{-3}} \right) \left( \frac{\tau_{\nuh_2}}{10^{-2}~{\rm s}} \right)^{1/2}, \end{eqnarray} where $a_j$ and $T_j$ (for $j=$bef and aft) are the scale factor and the photon temperature of the universe at time $j$, and we supposed $g_\ast=g_{\ast{\rm S}}$ and that the $g_\ast$ and $g_{\ast{\rm S}}$ values do not change between the temperatures of $T_{\rm bef}$ and $T_{\rm aft}$. We note that a large dilution factor is realized only for $M_{\nuh_2}\lesssim 100$ GeV. If the mass is much larger than the energy scale of the EW phase transition, the freeze-out $\nuh_2$ abundance is small because of the Boltzmann suppression factor. Furthermore, the lifetime should not be longer than ${\mathcal O}(10^{-2})$ s since BBN is significantly affected by nonthermal reactions of hadronic particles generated at the $\nuh_2$ decay if the lifetime is longer \cite{Kawasaki:2004qu}. From this equation, we find that the dilution factor, that equals the entropy enhancement factor, is about a factor of 100 at maximum. This maximum factor is $\sim$3 times smaller than the necessary factor of 300. Some other mechanism of the dilution is, therefore, needed for the light sterile neutrino to destroy some moderate fraction of primordial $^7$Be successfully. Possible mechanisms include dilutions by massive particles other than $\nuh_2$ decaying into active neutrinos, $\phi \rightarrow \nu \bar{\nu}$ \cite{Hannestad:2004px} or photons $\phi \rightarrow n\gamma$ ($n \geq 2$) \cite{Ichikawa:2005vw}. In Figs. \ref{fig1}, \ref{fig2}, \ref{fig3}, and \ref{fig1_old}, parameter regions of thermal freeze-out $\nuh$ abundances are shown by shaded regions. A possible dilution of the sterile neutrino by a factor of 100 is taken into account. The freeze-out abundances without dilutions are higher than the figure domains, and therefore not seen. The lower boundaries of these regions correspond to the abundances diluted by a factor of 100. Sudden drops of the boundaries at $\tau_\nuh =10^4-10^5$ s result from the decrease in the massless DOF in terms of entropy (Fig. \ref{fig_relic}). It is clear that the parameter regions for the primordial $^7$Li reduction are lower than the regions of freeze-out abundances. Therefore, a dilution of the sterile neutrino is necessary for the $^7$Li reduction to work. \subsection{mixing with muon and tauon neutrinos}\label{sec7c} If either muon or tauon neutrino predominantly couples to $\nuh$ and couplings of other charged leptons are negligible as an extreme case opposite to the case studied in this paper, effects of the sterile neutrino decay are changed. For example, we take the case of $M_\nuh=14$ MeV. The ratio of energy injections in the forms of $e^\pm$ and $\nu$ is $R(\nu, e)=3.20$ for the coupling to $\nu_e$, while it is $R(\nu, \alpha)=13.0$ for the coupling to $\nu_\alpha$ ($\alpha=\mu$ or $\tau$) [Eq. (\ref{eq_zeta_ratio})]. The muon or tauon type mixings, therefore, result in a large energy fraction of neutrino emitted at the decay. An $e^\pm$ injection decreases the $\eta$ and $N_{\rm eff}$ values and induces nonthermal nucleosynthesis, while a neutrino injection increases the $N_{\rm eff}$ value. Therefore, a sterile neutrino that mixes only with $\nu_\alpha$ has small effects on the primordial light element abundances and the $\eta$ value relative to that on $N_{\rm eff}$. \subsection{constraint from pion decay}\label{sec7d} We assume that the sterile neutrino has a mass $\sim 14$ MeV and a lifetime $\sim 10^5$ s (parameter value for the $^7$Li reduction), and that mixing angles of muon and tauon types, $\Theta_\mu$ and $\Theta_\tau$, respectively, can be neglected. The active-sterile mixing angle is then determined to be $\Theta ={\cal O}(10^{-3})$ [Eqs. (\ref{gamma_approx}) and (\ref{eq_a9})]. If the muon type mixing is sizable, we should take into account another constraint from low energy phenomena: the sterile neutrino can be produced by the decay of charged pions, e.g., $\pi^+ \to \mu^+ + \nuh$ or $\pi^+ \to e^+ + \nuh$. This channel has been searched for a long time and many experiments gave constraints on the active-sterile mixing angle. From Ref.~\cite{Kusenko:2004qc}, $\Theta^2_\mu$ should be smaller than $10^{-5}$. If the precision of those experiments can be improved by a few orders of magnitude, therefore, we may see a signal from pion decays, or exclude the possibility of primordial $^7$Li reduction by $\nuh$ suggested in this paper. Furthermore, if the muon type mixing is of the order of $\Theta_\mu \sim 10^{-6}$, the sterile neutrino might be detected by Super Kamiokande in future~\cite{Asaka:2012hc}. It is worth mentioning that a constraint from supernova SN1987A observation is rather strong~\cite{Dolgov:2000pj,Dolgov:2000jw}; that is roughly $\Theta^2 \lesssim {\cal O}(10^{-8})$ for any flavors.
14
3
1403.5995
The primordial lithium abundance inferred from spectroscopic observations of metal-poor stars is ∼3 times smaller than the theoretical prediction in the standard big bang nucleosynthesis (BBN) model. We assume a simple model composed of standard model particles and a sterile neutrino ν<SUB>H</SUB> with mass of O<SUB>(10) MeV which decays long after BBN. We then investigate cosmological effects of a sterile neutrino decay, and check if a sterile neutrino can reduce the primordial lithium abundance. We formulate the injection spectrum of nonthermal photon induced by the ν<SUB>H</SUB> decay. We take into account the generation of electrons and positrons, e<SUP>±</SUP>'s, and active neutrinos at the ν<SUB>H</SUB> decay, the primary photon production via the inverse Compton scattering of cosmic background radiation (CBR) by energetic e<SUP>±</SUP>, and electromagnetic cascade showers induced by the primary photons. The steady state injection spectrum is then derived as a function of the ν<SUB>H</SUB> mass and the photon temperature. The ν<SUB>H</SUB> decay produces energetic active neutrinos which are not thermalized, and e<SUP>±</SUP>'s which are thermalized. We then derive formulas relevant to the ν<SUB>H</SUB> decay rates and formulas for the baryon-to-photon ratio η and effective neutrino number Neff. The initial abundance, mass, and lifetime of ν<SUB>H</SUB> are taken as free parameters. We then consistently solve (1) the cosmic thermal history, (2) nonthermal nucleosynthesis induced by the nonthermal photons, (3) the η value, and (4) the N<SUB>eff</SUB> value. We find that an effective Be7 destruction can occur only if the sterile neutrino decays at photon temperature T =O(1) keV. Amounts of energy injection at the ν<SUB>H</SUB> decay are constrained from limits on primordial D and Li<SUB>7</SUB> abundances, the N<SUB>eff</SUB> value, and the CBR energy spectrum. We find that Be7 is photodisintegrated and the Li problem is partially solved for the lifetime 10<SUP>4</SUP>-10<SUP>5</SUP> s and the mass ≳14 MeV. Be7 destruction by more than a factor of 3 is not possible because of an associated D overdestruction. In the parameter region, the η value is decreased slightly, while the N<SUB>eff</SUB> value is increased by a factor of ΔN<SUB>eff</SUB>≲1. In this study, errors in photodisintegration cross sections of Be7(γ ,α)He3 and Li7(γ ,α)H3 that have propagated through the literature are corrected, and new functions are derived based on recent nuclear experiments. It is found that the new photodisintegration rates are 2.3 to 2.5 times smaller than the old rates. The correct cross sections thus indicate significantly smaller efficiencies of Be7 and Li7 photodisintegration. Abundances of sterile neutrino necessary for the Li7 reduction are much smaller than thermal freeze-out abundances. The relic sterile neutrino, therefore, must be diluted between the freeze-out and BBN epochs by some mechanism.</SUB>
false
[ "sterile neutrino", "energetic active neutrinos", "nonthermal photon", "active neutrinos", "effective neutrino number", "photon", "standard model particles", "H</SUB", "mechanism.</SUB", "energetic e", "a sterile neutrino decay", "decay", "cosmic background radiation", "Abundances", "BBN", "photodisintegration cross sections", "recent nuclear experiments", "electromagnetic cascade showers", "Be7 destruction", "new functions" ]
7.636301
-0.79819
-1
688585
[ "Guervilly, Céline", "Hughes, David W.", "Jones, Chris A." ]
2014JFM...758..407G
[ "Large-scale vortices in rapidly rotating Rayleigh-Bénard convection" ]
105
[ "-", "-", "-" ]
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[ "astronomy", "physics", "earth science" ]
11
[ "Physics - Fluid Dynamics", "Astrophysics - Earth and Planetary Astrophysics", "Physics - Geophysics" ]
[ "1961hhs..book.....C", "1967PhFl...10.1417K", "1982JFM...125..505H", "1984JFM...146...21M", "1986JFM...167..503B", "1987JFM...178...73G", "1987PhyD...28..382F", "1988JFM...195..303M", "1989JGR....9418445C", "1991PhFlA...3..403L", "1993AnRFM..25..241H", "1994Chaos...4..123B", "1994JFM...273....1B", "1996JFM...322..243J", "1997DyAtO..26..159Y", "1999PhRvE..59.4091L", "2002JFM...458..191V", "2002JFM...470..115C", "2003ApJ...588.1183C", "2003JFM...475...41C", "2004PhRvE..70e6312S", "2005CRPhy...6..431M", "2005Natur.438..193H", "2005PhFl...17i5105M", "2006JFM...551..141S", "2007AN....328.1059C", "2008JFM...597..395Z", "2008JFM...598...81S", "2009PhRvE..80a5305S", "2010GApFD.104..481S", "2011AN....332..876M", "2011ApJ...742...34K", "2012AnRFM..44..427B", "2012GApFD.106..392J", "2012JFM...691..568K", "2012PhRvE..85a6313K", "2013E&PSL.371..212C", "2014PhFl...26i6605F", "2014PhRvL.112n4501R" ]
[ "10.1017/jfm.2014.542", "10.48550/arXiv.1403.7442" ]
1403
1403.7442_arXiv.txt
The presence of large-scale coherent structures in turbulent flows attracts much interest, particularly because of their relevance in geophysics and astrophysics; understanding their formation is thus an important problem in fluid dynamics. In two-dimensional (2D) turbulence, in which vertical motions are assumed to be suppressed owing to strong stratification, fast rotation, or small vertical to horizontal scale ratio, the kinetic energy and the enstrophy (mean-square vorticity) are conserved quantities in the inviscid limit. This implies a downscale enstrophy cascade and an upscale energy cascade \citep{Kra67}, which can lead to the formation of coherent vortices \citep[e.g.][]{mcwilliams1984}. By contrast, in three-dimensional (3D) isotropic turbulence, enstrophy is not a conserved quantity, and the energy cascade is expected to be downscale. Nevertheless, the constraint imposed by rapid rotation might lead to 2D-like dynamics in a 3D flow on timescales longer than the rotation period. Notable examples of the formation of coherent vortices in a 3D system subject to rotation have been observed in experiments on grid-generated turbulence \citep[e.g.][]{Hopfinger1982, Staplehurst2008}. Recently, the presence of large-scale vortices (LSV) has been reported in numerical simulations of rotating convection in Cartesian geometry \citep{Chan07}, where the fluid is heated from below and confined between two horizontal planes. In this system, with buoyant vertical motions, the flow is necessarily $z$-dependent. The formation of LSV in convective layers still remains to be fully understood. Using numerical models of rotating compressible thermal convection in a local $f$-plane model, \citet{Chan07} and \citet{Chan13} report the emergence of long-lived, large-scale (\ie domain size) vortices for sufficiently large rotation rates. For moderate Rossby number ($\Ro$, the ratio of the rotation period to the typical convective turnover time), of the order of $0.1$, these LSV are cyclonic and associated with regions of lower temperature relative to their surroundings. Note that a vortex is defined as cyclonic (anticyclonic) when its vorticity in the rotating reference frame has the same (opposite) sign as the externally applied rotation. At lower Rossby numbers, \citeauthor{Chan07} and \citeauthor{Chan13} observe a large-scale warmer anticyclone accompanied by a smaller and weaker cyclone. Using a similar numerical set-up, \citet{Kap11} find that the LSV are excited provided that the Reynolds number ($\Rey$, the ratio of the viscous diffusion time to the convective turnover time) is sufficiently large. The vortices span the entire vertical extent of the box and are roughly aligned with the rotation axis. \citet{Man11} attribute the formation of these structures to a mean-field hydrodynamical instability that requires a sufficient scale separation between the convective eddies and the smallest horizontal wavenumber permitted in the computational domain. They find that increasing the box size leads to an increase of the horizontal extent of the structures, so that the LSV always fill roughly half of the horizontal domain. Large-scale structures have also been described in the work of \citet{Jul12}, who employ a set of reduced equations in a local Cartesian box describing Boussinesq convection in the limit of small Rossby number. In their model, the flow is locally in geostrophic balance at leading order $1/\Ro$, but thermally driven vertical flows exist at sufficiently small horizontal scales. When the thermal forcing is sufficiently large, \citeauthor{Jul12} observe the formation of a depth-invariant box-size flow, which becomes organised into a cyclone and anticyclone of similar strength. Using the same numerical model, \citet{Rubio13} show that the generation of these depth-invariant LSV involves the interactions of small-scale, depth-dependent convective eddies, which are made more coherent by the action of depth-invariant vortices. Interestingly, \citet{Jul12} find that the presence of LSV tends to increase the efficiency of the heat transfer through the system. Fully 3D Boussinesq convection in the presence of rotation has been extensively studied, particularly in cylindrical and spherical geometries with applications to the global dynamics of planetary interiors \citep[e.g.][]{Busse1994,Chr02}. In spherical geometry, the curved boundaries have an important effect on large-scale structures; in flows with large Reynolds numbers and low Rossby numbers they are, notably, responsible for the formation of zonal flows of amplitude large compared with the typical convective velocity \citep[e.g.][]{Hei05}. In simulations of rotating convection in spherical geometry, the formation of vortices at scales larger than the typical convective size has not been observed. \citet{Jul12} conjecture that, in a Cartesian domain, the size of the LSV is limited only by the domain size, so that if the upscale energy transfer were allowed to continue, the LSV would eventually feel the latitudinal variation of the Coriolis parameter. In this case, it is argued that the large-scale dynamics would become organised into zonal flows. That said, it is worth highlighting the occurrence of planetary polar vortices --- most strikingly those of Saturn --- the dynamics of which may be related to the dynamics of LSV in plane layer models. In numerical modelling, computational resources limit the values of parameters such as the Reynolds and Rossby numbers; this is even more pronounced in global spherical models compared with those in Cartesian geometry. Consequently, studies that aim to determine transitions between different convection regimes across a wide parameter range preferentially employ the rotating Rayleigh-B\'enard (RRB) configuration, in which a Boussinesq fluid contained between two horizontal planes rotates uniformly about an axis aligned with the direction of gravity \citep[e.g.][]{Julien1996, Vorobieff2002}. To our knowledge, among the previous studies of RRB convection conducted in the low Rossby number regime \citep{King2012, SchTil09, Stellmach2004}, the formation of box-size, vertically aligned vortices is addressed only in the contemporaneous study of \citet{Favier2014}. There are two possible explanations for the absence of LSV in most of the previous studies. One stems from the choice of boundary conditions, especially for the velocity. In the compressible convection models mentioned above, and also in the reduced Boussinesq model of \citet{Jul12}, stress-free boundary conditions are employed; often though, no-slip boundary conditions are adopted in RRB convection models \citep[e.g.][]{SchTil10, King2012}. Another plausible explanation for the lack of LSV stems from the choice of aspect ratio of the computational domain. In RRB simulations, the aspect ratio is usually taken equal to unity or smaller, whereas in the compressible convection studies, the aspect ratio is usually about four. Simulations of RRB convection are often carried out in order to assess the efficiency of heat transfer, thereby allowing the determination of the transition between rapidly rotating and non-rotating convection. As observed by \citet{Jul12}, heat transfer can be affected by the presence of LSV. It is therefore important to identify the conditions required for the formation of the LSV in RRB convection, and to assess their impact on heat transfer. In this paper, we investigate in detail the emergence of large-scale, depth-invariant vortices in convective regions via a series of numerical simulations of RRB convection. Our objectives are threefold: (i) to determine the parameters governing the presence of LSV; (ii) to understand the mechanism by which they form; (iii) to assess how LSV affect the heat transfer in the system. The layout of the paper is as follows. The mathematical and numerical formulation of the problem is contained in \S\,\ref{sec:MF}. The formation, maintenance and influence of the LSV are described in \S\,\ref{sec:LSV}. The spatial structure of the large-scale vortices, which always consist of a concentrated cyclone and a more dilute anticyclone, is discussed in \S\,\ref{sec:structure}, the domain of existence in parameter space in \S\,\ref{sec:domain}, and the reasons for the cyclonic/anticyclonic asymmetry in \S\,\ref{sec:cyclonic}. In \S\,\ref{sec:filter}, we establish how energy is transferred to the large scales. Finally, in \S\,\ref{sec:heat}, we discuss how the LSV affect the heat transfer in the system. A concluding discussion is contained in \S\,\ref{sec:discussion}.
\label{sec:discussion} We have presented simulations of rotating Rayleigh-B\'enard (RRB) convection that demonstrate the emergence of long-lived, large-scale vortices (LSV). These LSV consist of a patch of strong cyclonic vorticity surrounded by a region of weaker anticyclonic vorticity, both aligned with the rotation axis, which appear at the box size and are nearly depth-independent. With stress-free top and bottom boundaries, for the Ekman numbers considered here \mbox{($\Ek=10^{-4}$ -- $5\times 10^{-6}$)} and depending on the aspect ratio, the kinetic energy of the horizontal flow can be as much as ten times greater than that of the vertical flow, which is driven directly by buoyancy. LSV are observed when the Reynolds number based on the rms vertical velocity exceeds $100 \textrm{ -- } 300$, the threshold value being dependent on the box aspect ratio but independent of the Ekman number. This corresponds to Rayleigh numbers only about three times that at the onset of convection. The amplitude of the large-scale flow starts to decline once the thermal input is strong enough to allow a relaxation of the rotational constraint. Quantitatively, this decay of the LSV occurs for a local Rossby number based on the convective velocity, $\Ro_z^l$, of approximately $0.15$. Moreover, if the two conditions (i) \mbox{$\Rey_z > 100 \textrm{ -- } 300$} and (ii) \mbox{$\Ro_z^l \lesssim 0.15$} are met, we always observe a transfer of energy to the large horizontal scale, even for modest scale separation between the horizontal convective eddies and the horizontal extent of the domain (a factor four between the two is the smallest scale separation we considered). We tested the cyclone/anticyclone asymmetry of the LSV by artificially inverting the sign of the vorticity at a given time; the large-scale anticyclone subsequently disintegrates into smaller vortices, and the cyclone/anticyclone asymmetry at large scales is established relatively rapidly, after about $100$ rotation timescales. To gain some insight into the mechanism of the formation of the LSV, we performed a series of filtered simulations, in which spectral coefficients of given horizontal and vertical wavenumbers, $k_x$, $k_y$ and $k_z$, are artificially suppressed during the time integration. The filtered simulations suggest that the LSV (corresponding to $(k_x,k_y,k_z)=(1,1,0)$ in spectral space) are produced by the nonlinear interactions of small-scale $z$-dependent convective motions. Moreover, the presence of the spectral range between $(k_x,k_y)=(1,1)$ and the typical horizontal wavenumber of the convective structures is not required to sustain the LSV. As mentioned above, the amplitude of the LSV declines if the convection is not strongly influenced by rotation, in which case the convective structures are less anisotropic. To interact coherently, the convective motions must therefore present a significant anisotropy between their vertical and horizontal extents, \ie they must be significantly affected by rotation. In our study, the smallest compensated Rayleigh number, $\tRa$, at which LSV appear for different Ekman numbers corresponds to the transition from cellular convection ($\tRa \lesssim 20$) to the thermal plumes ($\tRa \gtrsim 20$) measured in the study of \citet{Jul12}, which is based on a reduced model of Boussinesq convection valid in the small Rossby number limit. Thermal plumes originate from a buoyancy instability in the thermal boundary layers. Vortex stretching within the plumes ejected from the boundaries yields an axial vorticity distribution that is skewed towards positive values near the top and bottom boundaries \citep{Chen1989,Julien1996}. Since the axial vorticity has no horizontal average, anticyclonic convective structures are necessarily also present, but they are less compact and have weaker vorticity. The formation of intense cyclonic thermal plumes near the boundaries could explain the predominance of the large-scale cyclonic circulation. Assuming that two cyclonic plumes form from the thermal boundary layer at sufficiently small distance, they would start to drift horizontally around one another \citep[e.g.][]{Boubnov1986,Hopfinger1993}; the conditions for the merger of two like-signed vortices depend notably on their separation distance, radius, and vorticity, and are the subject of an abundant literature on vortex dynamics \citep[e.g.][]{Griffiths1987, Melander1988, Cerretelli2003, Meunier2005}. The patch of cyclonic vorticity they create will be deformed by the background shear created by nearby individual vortices. In return, the deformed cyclonic patch tends to attract nearby cyclones and repel anticyclones \citep{Yasuda1997}. Merging of anticyclonic structures can also occur, but would be less likely as vortices of intense strength are more likely to merge. Since the horizontal boundaries are periodic, the repulsion of anticyclones by the large-scale cyclone would tend to group the anticyclones in the surrounding area, and so establish the weak anticyclonic circulation. In this scenario, the underlying asymmetry between large-scale cyclonic and anticyclonic circulation arises therefore through the formation of thermal plumes, which builds up a population of strong narrow cyclonic vortices, and the interactions between like-signed vortices then lead to the formation of one large cyclonic vortex by absorbing this available population of strong narrow cyclones. A potential weakness of this explanation for LSV formation is that the initial population of narrow cyclones is mainly located near the horizontal boundaries, whereas the LSV span the entire vertical extent of the domain. The clustering of cyclonic vorticity is described here as a two-dimensional process, but the conditions for interaction and merger of three-dimensional vortices have also been studied in detail in the literature \citep[e.g.][]{Ozuugurlu2008}. As observed during the time evolution of the kinetic energy of the large horizontal scale, the process of formation is slow, and occurs over thousands of rotation timescales; the large-scale flow eventually saturates when the viscous dissipation can balance the clustering of the convective eddies. In their reduced model, \citet{Jul12} observe the thermal plume regime for $\tRa \gtrsim 20$, but they report the formation of LSV for larger Rayleigh number, namely $\tRa=100$. In their study, the large-scale depth-invariant mode consists of a cyclone/anticyclone pair of similar vorticity. Towards the low Rossby number limit, the asymmetry between cyclonic and anticyclonic thermal plumes tends to vanish \citep{Vorobieff2002,Sprague2006}, and we indeed expect that the process of clustering of like-sign vorticity plumes would produce both large-scale cyclonic and anticyclonic circulation of equal strength. The simple scenario we propose for the formation of the large-scale cyclone is in agreement with the result of our filtered simulations, but remains to some degree speculative. This proposed picture could work in conjunction with the instability of large-scale anticyclonic regions for which the total vorticity ($\omega_z + 2 \Omega$) is small \citep{Lesieur91}. The reduction of the rotational constraint on the convection in large-scale anticyclonic regions yielding a possible local increase of $\Ro_z^l$ above the threshold value of $0.15$ could also contribute to the preference for cyclonic LSV. To confirm the proposed scenario of the formation of LSV, it would be interesting to study the interactions of a small number of convective structures in isolation, together with the effect of artificially added large patches of vorticity. Such studies are beyond the scope of the present paper, but could be addressed in future work. Despite a number of fully 3D models of rotating Boussinesq convection in Cartesian boxes, this work is one of the first to report the formation of box-size vortices in this system \citep[see also][]{Favier2014}. The independent work of \citet{Favier2014} has been carried out using a model of RRB convection with the same boundary conditions as ours, and their results are in agreement concerning the domain of existence of the LSV, the nature of the energy transfer to large scales, and the asymmetry between cyclones and anticyclones. An interesting difference between the two studies is that \citeauthor{Favier2014} use a computational domain of larger aspect ratio for similar Ekman numbers (for instance $\lambda=4$ for $\Ek=10^{-5}$), so they were able to achieve larger scale separation between the box size and the convective scales. In this case, they observe that several coherent cyclonic vortices coexist initially, and that these merge when two of these cyclones become close together, with eventually only one box-size cyclone remaining. Some of the earlier numerical studies of RRB convection have been carried out in the same parameter regime for which we identified the existence of LSV. Most of these were interested in measuring the heat flux in order to deduce scaling laws and thus identify transitions between the various convection regimes. As shown in \S\,\ref{sec:heat}, the presence of LSV markedly disturbs the convection by inhibiting the mixing in the core of the cyclone, yielding a reduction of the Nusselt number compared with its value when convection sets in, of about $5$ to $8$\%, depending on the aspect ratio. However, when viewed over several decades of the input parameters, this effect on the Nusselt number is not particularly noticeable on the scaling laws calculated with different aspect ratios. LSV could therefore be present in these earlier studies but not reported because of their minor influence on the scaling laws of the heat flux. The choice of boundary conditions is probably an important factor for the formation of LSV. No-slip boundary conditions for the velocity would tend to increase the viscous damping in the boundary layers compared with stress-free conditions, thereby reducing the amplitude of the horizontal flows. The absence of LSV in a number of experimental studies conducted in the range of parameters where LSV might be expected \citep[e.g.][]{Boubnov1986,Zhong09,King2012b} possibly suggests a destructive effect of no-slip boundaries on these structures. However, the comparison between simulations and experiments is not entirely straightforward because the main difference lies not only in the top and bottom boundary conditions, but also in the presence of side walls, which are known to influence convection in some cases \citep{Liu99}. Furthermore, in order to observe LSV in fluids with low viscosity, it is necessary to run experiments for a long time since the saturation depends on the viscous dissipation at large scales. Finally, changing the boundary conditions for the temperature to fixed flux rather than fixed temperature may also affect the presence of LSV, although the effect in this case is more difficult to predict. The effect of the boundary conditions therefore remains an interesting open question, which we propose to investigate in subsequent work.
14
3
1403.7442
Using numerical simulations of rapidly rotating Boussinesq convection in a Cartesian box, we study the formation of long-lived, large-scale, depth-invariant coherent structures. These structures, which consist of concentrated cyclones, grow to the horizontal size of the box, with velocities significantly larger than the convective motions. We vary the rotation rate, the thermal driving and the aspect ratio in order to determine the domain of existence of these large-scale vortices (LSV). We find that two conditions are required for their formation. First, the Rayleigh number, a meaure of the thermal driving, must be several times its value at the linear onset of convection; this corresponds to Reynolds numbers, based on the convective velocity and the box depth, $\gtrsim 100$. Second, the rotational constraint on the convective structures must be strong. This requires that the local Rossby number, based on the convective velocity and the horizontal convective scale, $\lesssim 0.15$. Simulations in which certain wavenumbers are artificially suppressed in spectral space suggest that the LSV are produced by the interactions of small-scale, depth-dependent convective motions. The presence of LSV significantly reduces the efficiency of the convective heat transport.
false
[ "the horizontal convective scale", "LSV", "velocities", "concentrated cyclones", "the convective heat transport", "the convective structures", "spectral space", "Reynolds numbers", "the convective velocity", "the convective motions", "certain wavenumbers", "Boussinesq convection", "small-scale, depth-dependent convective motions", "numerical simulations", "convection", "the box depth", "these large-scale vortices", "$\\lesssim 0.15$. Simulations", "Reynolds", "existence" ]
11.018406
14.324198
2
693817
[ "Hamaguchi, Koichi", "Liew, Seng Pei", "Moroi, Takeo", "Yamamoto, Yasuhiro" ]
2014JHEP...05..086H
[ "Isospin-violating dark matter with colored mediators" ]
27
[ "Department of Physics, University of Tokyo; Kavli Institute for the Physics and Mathematics of the Universe (WPI), University of Tokyo", "Department of Physics, University of Tokyo", "Department of Physics, University of Tokyo; Kavli Institute for the Physics and Mathematics of the Universe (WPI), University of Tokyo", "Departamento de Física Teórica y del Cosmos, Universidad de Granada" ]
[ "2014JCAP...08..011Z", "2014JHEP...06..169G", "2015JHEP...01..037B", "2015JHEP...04..054D", "2015JHEP...04..175L", "2015PhRvD..91c5009H", "2015PhRvD..91i2004S", "2015arXiv150301780M", "2016PhRvD..94c5012H", "2017JHEP...06..082L", "2017JHEP...10..059K", "2017JHEP...11..024L", "2017JHEP...11..196C", "2017PhRvD..95e5015Y", "2018JHEP...07..054D", "2018PhRvD..97e6004K", "2019JCAP...04..041Y", "2020PDU....2700371A", "2020PhRvD.101c5016L", "2020PhRvD.102h2001A", "2021JCAP...07..012B", "2022EPJC...82.1132A", "2022JCAP...06..026B", "2022JHEP...10..014B", "2022PhRvD.106g5010L", "2022arXiv220200996B", "2023PhRvD.108c1501C" ]
[ "astronomy", "physics" ]
7
[ "Phenomenological Models", "Hadronic Colliders", "High Energy Physics - Phenomenology", "Astrophysics - Cosmology and Nongalactic Astrophysics", "High Energy Physics - Experiment" ]
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[ "10.1007/JHEP05(2014)086", "10.48550/arXiv.1403.0324" ]
1403
1403.0324_arXiv.txt
Dark matter (DM), which is expected to be responsible for about $27\ \%$ of the mass density of the present universe \cite{Ade:2013zuv}, is still a great mystery to the field of particle physics. Although various cosmological observations have confirmed the existence of DM, its particle-physics properties, such as the mass, strength of its interactions with Standard--Model (SM) particles, and so on, remain fully unknown. Various experiments have been performed to detect direct and indirect signals of DM~\cite{Bertone:2004pz}. In recent years, several direct detection experiments (DAMA/LIBRA~\cite{Bernabei:2010mq}, CoGeNT~\cite{Aalseth:2010vx,Aalseth:2011wp,Aalseth:2014eft}, CRESST~\cite{Angloher:2011uu} and the CDMS-II Si experiment~\cite{Agnese:2013rvf}) have found signals that may suggest the existence of light DM with mass around 10 GeV. On the other hand, experiments such as~XENON~\cite{Angle:2011th,Aprile:2012nq}, LUX~\cite{Akerib:2013tjd}, SIMPLE~\cite{Felizardo:2011uw}, CDMS~\cite{Akerib:2010pv,Ahmed:2010wy} and SuperCDMS~\cite{Agnese:2013jaa, Agnese:2014aze} have not found any excess of events that can be interpreted as signals from DM. In particular, the LUX experiment has probed the relevant region of parameters at the highest level of sensitivity and excluded most regions favored by the possible signals of light DM. It has been shown that it is difficult to accommodate positive signals of direct detection experiments and bounds from Xenon-based experiments by considering astrophysical alternatives (e.g., modified halo models) or varying assumptions about the Xenon scintillation efficiencies~\cite{Gresham:2013mua,DelNobile:2013gba,Fox:2013pia}. A scenario that still remains viable in reconciling some of these results is the isospin--violating DM~\cite{Kurylov:2003ra,Giuliani:2005my,Chang:2010yk,Kang:2010mh,Feng:2011vu,Frandsen:2013cna,Feng:2013vod}. As different types of nuclei are used in different direct detection experiments, isospin--dependent interactions may happen to interfere destructively for a certain type of nuclei, and thus suppress the DM--nucleus scattering cross section. As LUX experiment~\cite{Akerib:2013tjd} currently imposes the most stringent bound on DM, one necessarily considers DM that has negligible interaction with the Xenon nucleus. Recent studies after the LUX result~\cite{Gresham:2013mua,DelNobile:2013gba,Cirigliano:2013zta,Fox:2013pia} have shown that the isospin--violating DM is still compatible with one of positive signals, those of the CDMS-II Si experiment. More recently, SuperCDMS Collaboration reported their first result for the WIMP search using their background rejection capabilities~\cite{Agnese:2014aze}. As we shall see, the isospin--violating DM scenario is severely constrained also by SuperCDMS, but there is still a viable region of parameter space. In this paper, we study a minimal extension of the SM with isospin--violating DM, assuming that the isospin--violating interaction of the DM is mediated by colored particles.\footnote{For recent studies on DM models with colored mediators, see, e.g., Refs.~\cite{Chang:2013oia,An:2013xka,Bai:2013iqa,DiFranzo:2013vra,Papucci:2014iwa}.} We investigate the phenomenology of the model, including collider searches as well as flavor and CP physics, paying particular attention to the parameter region which is consistent with CDMS--Si, LUX and SuperCDMS results. We show that a minimal viable model includes scalar DM and new colored vector-like fermions with masses of $O(1)$ TeV as mediators. The colored vector--like fermions can be tested at the 14 TeV LHC. On the other hand, the flavor and CP constraints severely restrict the parameters of the model. We also show that fermionic DM models with colored scalar mediators are disfavored by the LHC constraints. The remaining of the paper is organized as follows. In Section~\ref{s:eff}, we study effective operators involving SM and DM fields that reproduce the direct detection experimental results. We then study bounds on these operators from collider search and indirect detection. In Section~\ref{s:main}, we introduce a simple model involving only DM and colored mediators as new particles, that can reproduce the effective operators studied in Section~\ref{s:eff}. We study the current bound (8 TeV LHC) on the colored mediators and their prospects of discovery for 14 TeV LHC. In Section~\ref{SecFlavor}, we examine flavor and CP constraints on this model. Section~\ref{sec:conclusion} is devoted to conclusions. In Appendix~\ref{s:maj}, we briefly discuss models of isospin--violating fermionic DM mediated by colored scalars.
\label{sec:conclusion} In this paper, we have studied isospin--violating light DM that can explain the possible CDMS-Si signal of light DM, while avoiding the constraints by recent LUX and SuperCDMS experiments. In particular, we considered isospin--violating light DM models with colored mediators. We have shown that a minimal viable model includes scalar DM and new colored vector-like fermions with masses of $O(1)$ TeV as mediators. We investigated the collider searches, flavor and CP phenomenology. The masses of colored mediators are constrained by the 8 TeV LHC results as $M\gtrsim 1-1.5$ TeV ($1-1.1$ TeV) for real (complex) scalar DM. The 14 TeV LHC may cover a large region of the remaining parameter space. We have also studied flavor and CP constraints on the colored-mediator model for the isospin--violating DM. In such a model, the interaction of quarks with colored mediator and DM should be sizable, which results in large radiative correction to flavor and CP observables. We have studied the effects on the quark masses (in particular, those of up- and down-quarks), EDM of neutron, and the $K$-$\bar{K}$ mixing parameters. Radiative corrections to the SM Yukawa couplings from the DM sector are extremely large, and hence fine--tunings are unavoidable. Flavor and CP violating observables also impose severe constraints on the present scenario.
14
3
1403.0324
In light of positive signals reported by the CDMS-II Si experiment and the recent results of the LUX and SuperCDMS experiments, we study isospin-violating dark matter scenarios assuming that the interaction of the dark matter is mediated by colored particles. We investigate the phenomenology of the model, including collider searches, flavor and CP phenomenology. A minimal possible scenario includes scalar dark matter and new vector-like colored fermions with masses of O(1) TeV as mediators. Such a scenario may be probed at the 14 TeV LHC, while flavor and CP constraints are stringent and severe tuning in the couplings is unavoidable. We also found that, as an explanation of the CDMS-II Si signal, isospin-violating fermionic dark matter models with colored scalar mediators are disfavored by the LHC constraints.
false
[ "scalar dark matter", "CP constraints", "colored scalar mediators", "TeV LHC", "CP phenomenology", "isospin-violating fermionic dark matter models", "colored particles", "isospin-violating dark matter scenarios", "SuperCDMS experiments", "mediators", "flavor", "stringent and severe tuning", "the LHC constraints", "the dark matter", "new vector-like colored fermions", "O(1", "masses", "positive signals", "searches", "LUX" ]
8.328863
-1.915541
54
812698
[ "Biern, Sang Gyu", "Gong, Jinn-Ouk", "Jeong, Donghui" ]
2014PhRvD..89j3523B
[ "Nonlinear matter bispectrum in general relativity" ]
6
[ "Department of Physics, Seoul National University, Seoul 151-747, Korea", "Asia Pacific Center for Theoretical Physics, Pohang 790-784, Korea; Department of Physics, Postech, Pohang 790-784, Korea", "Department of Physics and Astronomy, Johns Hopkins University, Baltimore, Maryland 21218, USA" ]
[ "2014PhRvD..90l3507Y", "2015CQGra..32q5019B", "2015arXiv150502972B", "2016JCAP...01..030V", "2016JCAP...07..017Y", "2017JCAP...03..034U" ]
[ "astronomy", "physics" ]
6
[ "98.80.-k", "04.20.-q", "98.65.Dx", "Cosmology", "Classical general relativity", "Superclusters", "large-scale structure of the Universe", "Astrophysics - Cosmology and Nongalactic Astrophysics", "General Relativity and Quantum Cosmology" ]
[ "1980PhRvD..22.1882B", "1994ApJ...431..495J", "1997ApJ...487....1S", "2002PhR...367....1B", "2004PhRvD..69j4011N", "2006ApJ...651..619J", "2008GReGr..40.1997A", "2009PhRvD..80h3514Y", "2010CQGra..27l4009B", "2011ApJ...727...22J", "2011ApJS..192...18K", "2011JCAP...07..025K", "2011MNRAS.418.1707B", "2011PhRvD..84d3516C", "2011PhRvD..84f3505B", "2012ApJ...752...50H", "2012PhRvD..85b3504J", "2013A&A...557A..54D", "2013ApJS..208...19H", "2013ApJS..208...20B", "2013PhR...530...87W", "2014A&A...571A...1P", "2014A&A...571A..16P", "2014JCAP...02..006A", "2014JLTP..176..733M", "2014MNRAS.441...24A", "2015MNRAS.446..677R" ]
[ "10.1103/PhysRevD.89.103523", "10.48550/arXiv.1403.0438" ]
1403
1403.0438_arXiv.txt
\label{sec:intro} Recent advances in cosmology has been greatly spurred by precise cosmological observations. The accurate measurements of the temperature anisotropies and polarizations in the cosmic microwave background (CMB) by the Wilkinson Microwave Anisotropy Probe (WMAP) have opened the era of precision cosmology~\cite{wmap}, and with the most recent PLANCK data we can constrain the cosmological parameters with less than $\calO(1)$ percent error~\cite{planck}. With the planned experiments such as PIXIE~\cite{PIXIE}, PRISM~\cite{PRISM}, and LiteBIRD~\cite{LiteBIRD} to mention a few, it is guaranteed that we continue our success in the CMB observations and that we can constrain the cosmological parameters further and can obtain more information on the early universe as well. Large scale structure (LSS) of the universe is yet another powerful cosmological probe, and its importance has ever been increasing with galaxy surveys such as SDSS~\cite{Anderson:2013zyy}, WiggleZ~\cite{Blake:2011en} and VIPERS~\cite{delaTorre:2013rpa}. The LSS observations can provide the measurement of geometrical distances, growth of structures, and shape of primordial correlation functions. These lower redshift information combined with the CMB data can break down the degeneracies among cosmological parameters that yields better constraints than CMB alone~\cite{planck}. Furthermore, the full three-dimensional information with a huge redshift coverage available for the LSS observations naturally yields measurement of properties of dark energy, neutrino properties as well as physics of the early universe. A number of future observations such as HETDEX~\cite{HETDEX}, MS-DESI~\cite{MSDESI}, LSST~\cite{LSST} and Euclid~\cite{Euclid} are proposed to observe LSS with improved accuracy in near future. Provided that unprecedentedly accurate data will be soon available in both CMB and LSS, our theoretical endeavour should also meet the observational precision. This introduces, however, a number of interesting and important questions to be addressed, especially for LSS: \begin{itemize} \item Non-linearity: With increasing observational accuracy, we can probe the signal beyond the two-point correlation function in CMB and LSS. The higher-order correlation functions are the signature of non-linearities. Searching for the primordial non-Gaussianity~\cite{nGreviews} is a prime example. The current best constraint from PLANCK is consistent with that the primordial fluctuations follow the Gaussian statistics with the local non-linearity parameter $\fnl = 2.7 \pm 5.8$ at $2\sigma$ confidence level. Non-linearity is more prominent in LSS: gravitational instability amplifies the density fluctuations to form non-linear structures such as galaxies and clusters of galaxies. As a result, the non-linearities deviates the matter power spectrum from the linear theory predictions~\cite{nlpowerspectrum,Jeong:2010ag}, and generates large higher-order correlation functions such as bispectrum and trispectrum. Accurate modeling of non-linearities is, therefore, the key requirement of exploiting the LSS data at the accuracy level similar to the CMB. \item Relevance of general relativity: Most studies on LSS in the past have been done in the context of the Newtonian gravity~\cite{Bernardeau:2001qr}, which works fine in the small scale, sub-horizon limit. In order to achieve robust measurements of dark energy properties, for example from Baryon Acoustic Oscillations (See~\cite{Weinberg:2012es} for a recent review), planned future LSS surveys will probe larger and larger volume, and access the scales comparable to the horizon. Modeling the LSS observables on those large scales demands that we work in the fully general relativistic context. The first question that must be addressed is whether the purely relativistic effects are large enough to be detected or not. Furthermore, attempted modifications to general relativity (to explain the recent cosmic acceleration) mostly show up on such very large scales. Thus LSS is a perfect playground to test modified theories of gravity. \item Gauge: As we should resort to general relativity, at least in principle, to study LSS properly, it is crucial to clarify which `gauge' we are using to interpret the data from LSS surveys. Different gauges are mathematically equivalent, but it does not mean that physical clarity is also equally shared. In particular, in the small scale limit the `density contrasts' $\delta\equiv T^0{}_0/\overline{T}^0{}_0 -1$ in almost all popular gauges are equivalent to the Newtonian density contrast~\cite{Hwang:2012bi}, but equivalence does not hold on large enough scales close to the horizon. Of course, by properly choosing the gauge that we interpret the data, the gauge ambiguity on large scales disappears to yield the gauge invariant expression for the observable such as the galaxy power spectrum~\cite{gaugedep}. \end{itemize} Bearing these in mind, we are encouraged to go beyond the two-point correlation function or power spectrum, and study the higher-order correlation functions arisen from the non-linearity in general relativity. In this article, we study the next-to-leading order non-linearities in the matter bispectrum in the comoving gauge. The non-linear matter power spectrum in the same gauge was computed in~\cite{Jeong:2010ag}. In the comoving gauge, the physical interpretation of the relativistic variables is transparent and the set of dynamical equations becomes particularly simple. Furthermore, the equations governing the dynamics of the density and velocity fields exactly coincide with the usual Newtonian hydrodynamic equations up to second order~\cite{Noh:2004bc}. Therefore, the leading order matter bispectrum, which results from correlating one second order density contrast to two linear order ones, in the full relativistic calculation must be the same as that of the Newtonian calculation, and the purely relativistic contributions appear from the third order. To obtain the self-consistent next-to-leading order non-linearities, we calculate the density contrast to the fourth-order. We compute the one-loop matter density and velocity bispectra, and confirm that the purely relativistic corrections are subdominant on cosmologically relevant scales. Going beyond the comoving gauge, we also calculate the leading order matter bispectrum from various other gauges to demonstrate the wild gauge dependence of the density and velocity bispectra. As in the case for the galaxy power spectrum, such a gauge dependence should go away when one calculate the `observable' quantities in each gauge. This article is organized as follows. In Section~\ref{sec:setup} we present the perturbation equations of a pressureless matter in the comoving gauge. In Section~\ref{sec:bispectrum} we give the fourth order solutions of the perturbation equations in terms of kernels, and compute the matter bispectrum including one-loop corrections. In Section~\ref{sec:result} we show the total bispectrum in particular configurations of interest. In Section~\ref{sec:gauge} we show gauge dependence of the leading bispectrum in general relativity for large scale study. We conclude in Section~\ref{sec:conclusion}.
\label{sec:conclusion} We have studied how the non-linearities in general relativistic affects the non-linear density and velocity bispectra. Using the full general relativistic formalism, we calculate one-loop bispectra of density and velocity fields in a flat, matter dominated universe. We have assumed that the initial density perturbation is perfectly Gaussian, so that the matter bispectrum comes solely from the non-linear dynamics. As we work in the comoving gauge, where the relativistic density and velocity fields coincide with those in the Newtonian theory, the pure relativistic corrections appear from third order. We have computed the non-linear bispectrum in the equilateral, folded and squeezed triangular configurations and have shown that the relativistic effects are appreciable only on the scale larger than the horizon. On small scales, the Newtonian one-loop corrections dominate the relativistic ones and even the tree contributions at $k\gtrsim0.2h$Mpc${}^{-1}$, indicating the non-linear evolution of the bispectrum due to gravitational instability. The general relativistic corrections appear to be dominant over the Newtonian ones when the longest wavemode is near comoving horizon scale. That is the reason why we see the domination of the relativistic corrections in the squeezed configurations on very large scales. We then have demonstrated the gauge dependence of the matter bispectrum by explicitly computing the tree-level matter bispectrum in four different gauges: comoving, synchronous, zero shear and uniform curvature gauges. Except for the synchronous gauge, whose time coordinate is comoving with the observer and thus the meaning of the density contrast differs from the other gauges even on the small scales, the matter bispectrum computed in the other gauges are the same on small scales. On large scales, the gauge dependence begin to show up more prominently and the matter bisepctra from all four gauges diverge from each other. This result is, again, consistent with the one-loop result in the comoving gauge that the general relativistic effects are only important near the horizon scales. The gauge dependence that we see near the horizon scale is the outcome of that the matter bispectrum itself is not a direct observable. When considering the observable quantities such as the bispectrum of weak-lensing shear and convergence field, or the galaxy bispectrum including all the relevant effects such as galaxy bias, light reflection, etc, the gauge dependence must disappear. In the case of the galaxy power spectrum, \cite{gaugedep} have shown that the observed galaxy power spectrum written in terms of the observed coordinate system is indeed gauge independent. Likewise, we surmise that calculating the bispectrum of observed quantities should resolve this gauge dependent ambiguities. Such a calculation requires extending the previous work done for the galaxy power spectrum to second order. \subsection*{Acknowledgements} SGB appreciates Jai-chan Hwang for suggesting helpful comments and supporting this research. We thank the Topical Research Program ``Theories and practices in large scale structure formation'', supported by the Asia Pacific Center for Theoretical Physics, while this work was under progress. JG acknowledges the Max-Planck-Gesellschaft, the Korea Ministry of Education, Science and Technology, Gyeongsangbuk-Do and Pohang City for the support of the Independent Junior Research Group at the Asia Pacific Center for Theoretical Physics. JG is also supported by a Starting Grant through the Basic Science Research Program of the National Research Foundation of Korea (2013R1A1A1006701). DJ is supported by DoE SC-0008108 and NASA NNX12AE86G, and acknowledges support from the John Templeton Foundation. \appendix
14
3
1403.0438
We show that the relativistic effects are negligibly small in the nonlinear density and velocity bispectra. Although the nonlinearities of the Einstein equation introduce additional nonlinear terms to the Newtonian fluid equations, the corrections to the bispectrum only show up on superhorizon scales. We show this with the next-to-leading order nonlinear bispectrum for a pressureless fluid in a flat Friedmann-Robertson-Walker background by calculating the density and velocity fields up to fourth order. We work in the comoving gauge, where the dynamics is identical to the Newtonian up to second order. We also discuss the leading order matter bispectrum in various gauges, and show yet another relativistic effect near horizon scales that the matter bispectrum strongly depends on the gauge choice.
false
[ "fourth order", "second order", "various gauges", "additional nonlinear terms", "bispectrum", "superhorizon scales", "the leading order matter", "fourth", "the nonlinear density and velocity bispectra", "second", "the Newtonian fluid equations", "Newtonian", "the gauge choice", "a pressureless fluid", "the density and velocity fields", "horizon", "the comoving gauge", "a flat Friedmann-Robertson-Walker background", "the relativistic effects", "the bispectrum" ]
11.67163
2.124399
-1
746694
[ "Generozov, A.", "Haiman, Z." ]
2014MNRAS.443L..64G
[ "Lyman edges in supermassive black hole binaries." ]
7
[ "Department of Astronomy, Columbia University, 550 West 120th Street, New York, NY 10027, USA", "Department of Astronomy, Columbia University, 550 West 120th Street, New York, NY 10027, USA" ]
[ "2014ApJ...789..140L", "2014arXiv1410.3563M", "2015MNRAS.446L..36F", "2017ApJ...834..200M", "2018MNRAS.474.2975D", "2022A&A...663A..99K", "2023arXiv231016896D" ]
[ "astronomy" ]
4
[ "accretion", "accretion discs", "black hole physics", "galaxies: active", "Astrophysics - High Energy Astrophysical Phenomena" ]
[ "1973A&A....24..337S", "1980Natur.287..307B", "1984AdSpR...3j.249K", "1989ApJ...342...64A", "1995ARA&A..33..581K", "1995ApJ...439..875H", "1996ApJ...467L..77A", "1997ApJ...484L..37H", "1999ApJS..123....3L", "1999PASP..111....1K", "2000ApJ...528..637B", "2003A&A...397..527S", "2003ApJ...582L..15K", "2003ApJ...593...69K", "2005ApJ...622L..93M", "2005SSRv..116..523F", "2006ApJ...646...49R", "2006MmSAI..77..733K", "2007PASJ...59..427H", "2008ApJ...672...83M", "2009ApJ...700.1952H", "2009CQGra..26i4032H", "2009MNRAS.393.1423C", "2009MNRAS.398.1392L", "2009astro2010S.235P", "2010ApJ...714..404T", "2011ApJ...741...57D", "2011Natur.477..431F", "2012A&A...545A.127R", "2012ApJ...749..118S", "2012ApJ...755....7D", "2012ApJ...755...51N", "2012ApJ...761...90G", "2012MNRAS.420..705T", "2012MNRAS.420..860S", "2012MNRAS.427.2660K", "2013ApJ...775...49S", "2013CQGra..30v0301B", "2013MNRAS.434.2275T", "2013MNRAS.436.2997D", "2014ApJ...783..134F", "2014ApJ...785..115R", "2014MNRAS.438.3024L", "2014PhRvD..89f4060G" ]
[ "10.1093/mnrasl/slu075", "10.48550/arXiv.1403.0002" ]
1403
1403.0002_arXiv.txt
\label{sec:introduction} Supermassive black holes (SMBHs) are present in the centres of most, if not all, nearby galaxies (see reviews by e.g. \citealt{KR95}; \citealt{FF05}). If two galaxies containing SMBHs merge, this should then result in the formation of a SMBHB (e.g. \citealt{Begel:Blan:Rees:1980}). Thus, given the hierarchical model for structure formation, in which galaxies are built up by mergers, one would naively expect SMBHBs to be quite common. Many candidates for binary BHs have been identified on kiloparsec scales, including two galaxies with spatially resolved active binary nuclei (\citealt{Komossa+2003,Fabbiano+2011}; see, e.g. the review by \citealt{Komossa:Rev06} and \citealt{Shen+:2013} and references therein). At parsec scales, however, there is only one clear example: a radio observation of a BH pair with a projected separation of $\sim$7 pc \citep{Rodriguez+:2006}. There remain no confirmed binary black holes at subparsec separations. The lack of observational evidence for binaries at small separations suggests that the SMBHs either remain inactive during the merger, or that they merge within a small fraction of a Hubble time and are consequently rare (\citealt{HKM09}). Another possibility is that the spectrum of a compact binary differs significantly from those of single-BH active galactic nuclei (AGN). A better understanding of the spectral energy distributions (SEDs) and lightcurves from circumbinary discs is necessary to determine whether binaries may therefore be missing from AGN surveys or catalogs \citep{Tanaka:2013}. Gravitational waves (GWs) from a merging SMBHB may be detected in the next decade by pulsar timing arrays (PTAs; \citealt{PTAs}). Identifying the gravitational wave source in EM bands would also have considerable payoffs for cosmology and astrophysics (e.g. \citealt{PhinneyWhitePaper}). Unfortunately, GWs yield limited precision on the sky position. For a PTA source, of order 10$^2$ (and perhaps as many as 10$^4$) plausible candidates may be present within the 3D measurement error box \citep{Tanaka:2012}. Concurrent EM observations would then be necessary to identify the GW source. Many different EM signatures have been proposed for SMBHBs. These include periodic luminosity variations commensurate with the orbital frequency of the source (\citealt{Haiman+2009} and references therein) and broad emission lines that are double-peaked and/or offset in frequency (\citealt{Shen+:2013} and the references therein). Additionally, the evacuation of a central cavity by the binary could lead to a spectrum that is a distinctively soft \citep{Milos:Phinney:2005,TanakaMenou:2010}, and has unusually weak broad optical emission lines, compared to typical AGN \citep{Tanaka:2012}. Other work \citep{Gultekin:Miller:2012, Roedig+:2014} describes spectral signatures of discs with partial cavities, which may show up in the SED as broad, shallow dips. We here propose that, in addition to the above signatures, the presence of a central cavity in a circumbinary disc could produce distinct absorption edges in the optical/UV (in particular, at the Lyman limit). This is analogous to the prominent Lyman break in galactic spectra a $\gsim$ few$\times10$ Myr after a starburst, when the composite emission is dominated by less massive stars with cooler atmospheres \citep{Leitherer+1999,Schaerer2003}. Physically, such an edge would be present if the disc is cold enough ($T_{\eff}\lsim20,000$ K) to have sufficient neutral (ground-state) hydrogen to absorb Lyman continuum photons. The disc should also be hot enough (conservatively $T_{\eff}\gsim10,000$ K); for yet cooler discs, the continuum emission redward of the Lyman limit may be obscured by metal absorption features. This Letter is organized as follows. In\,\S\ref{sec:discmodels}, we describe the details of our disc and emission models. In\,\S\ref{sec:results}, we show examples of spectra for binary BH discs, and compare these to those of single BH discs. In\,\S\ref{sec:minidiscs}, we discuss caveats, including emission from mini-discs around each of the individual BHs that could mask the Lyman edge. We summarize our main conclusions and the implications of this work in\,\S\ref{sec:conclusion}. \vspace{-\baselineskip}
\label{sec:conclusion} In this Letter, we have proposed that spectral edges, in particular at the Lyman limit, may be characteristic signatures of a circumbinary disc. Our conclusions can be summarized as follows. \begin{enumerate} \item If binary torques clear a cavity in the circumbinary disc, the disc spectrum may exhibit a sharp drop at the Lyman limit. This is because the hottest region (i.e. the inner edge) of the disc is cool enough to have neutral H, absorbing nearly all flux blueward of the Lyman limit. This occurs below a critical $\teff$ which generally lies in the range $\sim$10,000 K-20,000 K, depending on vertical gravity and other parameters. At lower temperatures, absorption from metals (i.e. C) may cause spectral edges redward of the Lyman limit. \item Observationally, AGN spectra only show Lyman edges due to absorption by intervening neutral gas (see \citealt{Antonucci+:1989}). The inner regions of a single-BH AGN disc are hotter than for binaries (Fig.~\ref{fig:temperature}), and can mask any edge produced in the outer disc, leaving only a small ``kink''(Fig.~\ref{fig:composite}). Such kinks are not seen observationally, and understanding what would smear them (e.g. general relativistic effects or winds) is an open theoretical problem. For an overview of the Lyman edge problem in AGN spectral modeling see, e.g. \citet{KolykhalovSunyaev1984} and \citet{KoratkarBlaes:99}. \item Neutral H in the binary's host galaxy, unrelated to the nuclear accretion disc itself, could cause a Lyman edge (as seen in a few AGNs). However, in the case of the disc, one could look for rotational broadening of the edge due to orbital motions in the disc, with velocities of order $10^4$ km/s. This may cause a $\simeq$10\% smearing on the edge. \item A portion of the binary parameter space could have a truncated circumbinary disc, with $\teff$ in the critical range for a prominent Lyman edge (Fig.~\ref{fig:param_space}). This parameter space partially overlaps with the expected typical parameters for individually resolvable PTA sources. These are very massive binaries, with $10^8 \Msun \lsim M \lsim 10^9 \Msun$, and separations ranging from 10's to 1000's of $R_g$, and mass ratios peaking at $q\sim1$ but with a long tail to lower values \citep{Sesana+2012}. \item Efficient fueling of the BHs inside the central cavity could mask the Lyman edge feature in the circumbinary disc spectrum. For example, persistent emission from hot mini-discs (see~\citealt{Farris+2013}) would obscure the Lyman edge. However, if the accretion flows onto the individual BHs are radiatively inefficient and/or intermittent \citep{Tanaka:2013}, the Lyman edge could remain visible, or appear periodically on the time-scale of the binary's orbit (which could be weeks to years; \citealt{HKM09}). \item The proposed Lyman edge signature could be used in combination with other proposed EM signatures to refine the search for SMBHBs. We have conducted a preliminary search for the Lyman edge feature in x-ray weak quasars discussed in \citealt{Brandt+:2000}. In particular we looked at FUSE and IUE spectra for the ten objects in their Table 2, but found no sign of any Lyman edge feature. \end{enumerate} If detected in an AGN, a prominent Lyman edge would tighten the case for the presence of a compact binary BH. \vspace{\baselineskip} We thank Shane Davis for sharing a modified version of \texttt{TLUSTY} and for technical help, as well as for providing an initial table of spectral models. We also thank Shane Davis, Omer Blaes, Ivan Hubeny, Jules Halpern, and Frits Paerls for insightful conversations. ZH acknowledges support from NASA grant NNX11AE05G. \footnotesize{
14
3
1403.0002
We propose a new spectral signature for supermassive black hole binaries (SMBHBs) with circumbinary gas discs: a sharp drop in flux bluewards of the Lyman limit. A prominent edge is produced if the gas dominating the emission in the Lyman continuum region of the spectrum is sufficiently cold (T ≲ 20 000 K) to contain significant neutral hydrogen. Circumbinary discs may be in this regime if the binary torques open a central cavity in the disc and clear most of the hot gas from the inner region, and if any residual UV emission from the individual BHs is either dim or intermittent. We model the vertical structure and spectra of circumbinary discs using the radiative transfer code TLUSTY, and identify the range of BH masses and binary separations producing a Lyman edge. We find that compact supermassive (M ≳ 10<SUP>8</SUP> M<SUB>⊙</SUB>) binaries with orbital periods of ∼0.1-10 yr, whose gravitational waves are expected to be detectable by pulsar timing arrays, could have prominent Lyman edges. Such strong spectral edge features are not typically present in AGN spectra and could serve as corroborating evidence for the presence of an SMBHB.
false
[ "prominent Lyman edges", "significant neutral hydrogen", "circumbinary gas discs", "Lyman", "continuum region", "Circumbinary discs", "circumbinary discs", "BH masses", "Such strong spectral edge features", "binary separations", "SMBHB", "pulsar timing", "supermassive black hole binaries", "AGN spectra", "BH", "a Lyman edge", "flux bluewards", "orbital periods", "the Lyman limit", "binaries" ]
7.937779
3.274493
39
482959
[ "Goodman, Michael L." ]
2014ApJ...785...87G
[ "Acceleration of Type 2 Spicules in the Solar Chromosphere. II. Viscous Braking and Upper Bounds on Coronal Energy Input" ]
8
[ "Advanced Technologies Group, West Virginia High Technology Consortium Foundation, 1000 Galliher Drive, Fairmont, WV 26554, USA" ]
[ "2014Sci...346A.315T", "2015LRSP...12....2L", "2015SSRv..190..103J", "2016IAUTA..29..278C", "2016SoPh..291.1129N", "2018A&A...616A..99K", "2019Sci...366..890S", "2020A&A...643A.140M" ]
[ "astronomy" ]
4
[ "magnetohydrodynamics: MHD", "stars: chromospheres", "stars: coronae", "Sun: chromosphere", "Sun: magnetic fields", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1961ApJ...134..347O", "1977ARA&A..15..363W", "1981ApJS...45..635V", "1984oup..book.....M", "1989ApJ...346.1010A", "1996ApJ...460.1019L", "1996ApJ...462..547L", "1998A&A...333.1053L", "2000ApJ...533..501G", "2004A&A...416.1159G", "2004ApJ...606..565B", "2005A&A...436..325L", "2006ASPC..358..448R", "2006ApJ...648L..67V", "2007ApJ...655..624D", "2007ApJ...666..541A", "2007ApJ...670L..61O", "2007PASJ...59S.655D", "2007Sci...318.1574D", "2008ApJ...672.1237L", "2008ApJ...679L.167L", "2009ApJ...702..129M", "2009ApJ...705..272R", "2009ApJ...706L..80M", "2009SSRv..144..197L", "2009SSRv..144..275D", "2010ApJ...708.1579C", "2010ApJ...719..469J", "2011ASPC..437..451S", "2011ApJ...736....9M", "2011NewA...16..296T", "2011Sci...331...55D", "2012A&A...541A..17S", "2012ApJ...748...77S", "2012ApJ...750L..25J", "2012ApJ...751...75G", "2012ApJ...752..108S", "2012ApJ...752L..12D", "2012ApJ...757..188G", "2012ApJ...757L..17D", "2012ApJ...759...18P", "2012JGRA..11712102K", "2012SSRv..169..181T", "2013A&ARv..21...66S" ]
[ "10.1088/0004-637X/785/2/87", "10.48550/arXiv.1403.2694" ]
1403
1403.2694_arXiv.txt
14
3
1403.2694
A magnetohydrodynamic model is used to determine conditions under which the Lorentz force accelerates plasma to type 2 spicule speeds in the chromosphere. The model generalizes a previous model to include a more realistic pre-spicule state, and the vertical viscous force. Two cases of acceleration under upper chromospheric conditions are considered. The magnetic field strength for these cases is &lt;=12.5 and 25 G. Plasma is accelerated to terminal vertical speeds of 66 and 78 km s<SUP>-1</SUP> in 100 s, compared with 124 and 397 km s<SUP>-1</SUP> for the case of zero viscosity. The flows are localized within horizontal diameters ~80 and 50 km. The total thermal energy generated by viscous dissipation is ~10 times larger than that due to Joule dissipation, but the magnitude of the total cooling due to rarefaction is &gt;~ this energy. Compressive heating dominates during the early phase of acceleration. The maximum energy injected into the corona by type 2 spicules, defined as the energy flux in the upper chromosphere, may largely balance total coronal energy losses in quiet regions, possibly also in coronal holes, but not in active regions. It is proposed that magnetic flux emergence in intergranular regions drives type 2 spicules.
false
[ "total coronal energy losses", "active regions", "quiet regions", "intergranular regions", "coronal holes", "terminal vertical speeds", "viscous dissipation", "upper chromospheric conditions", "Joule dissipation", "s", "The total thermal energy", "magnetic flux emergence", "the energy flux", "2 spicule speeds", "conditions", "G. Plasma", "The maximum energy", "the vertical viscous force", "acceleration", "type 2 spicules" ]
12.368176
15.35855
2
12372962
[ "Kanarek, G.", "Shara, M.", "Faherty, J.", "Zurek, D.", "Moffat, A." ]
2015MNRAS.452.2858K
[ "A near-infrared survey of the inner Galactic plane for Wolf-Rayet stars - III. New methods: faintest WR stars" ]
21
[ "Columbia University, 116th St &amp; Broadway, New York, NY 10027, USA", "American Museum of Natural History, 79th Street and Central Park West, New York, NY 10024, USA", "Department of Terrestrial Magnetism, Carnegie Institution of Washington, 5241 Broad Branch Road NW, Washington, DC 20015, USA", "American Museum of Natural History, 79th Street and Central Park West, New York, NY 10024, USA", "Départment de Physique, Université de Montréal, CP 6128 Succ. C-V, Montréal, QC H3C 3J7, Canada" ]
[ "2014AJ....147..115F", "2014AJ....148...34F", "2014ApJ...796..116N", "2015MNRAS.447.2322R", "2015arXiv150402477O", "2017IAUS..323...11F", "2017MNRAS.465..293K", "2017MNRAS.466..749M", "2017NewAR..79...49V", "2018MNRAS.473.2853R", "2018MNRAS.479.3759G", "2018MNRAS.480.2423I", "2019ApJ...885...68B", "2019MNRAS.488.1282W", "2020ApJ...898...74L", "2020MNRAS.495.1209R", "2022ApJ...937...46K", "2023MNRAS.525.3195S", "2024A&A...686A.118C", "2024ApJ...968...60K", "2024MNRAS.528.4657K" ]
[ "astronomy" ]
10
[ "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1867CRAS...65..292W", "1985ApJ...288..618R", "1987PASP...99..191S", "1993ASPC...52..246L", "1994PASP..106..250S", "1995A&A...293..889P", "1996AJ....112.2184G", "1997ApJ...486..420F", "1999AJ....118..390S", "2003A&A...408.1029K", "2004PASP..116..352V", "2004PASP..116..362C", "2005ApJ...619..931I", "2006A&A...458..453V", "2006AJ....131.1163S", "2006Ap&SS.304...97S", "2006MNRAS.372.1407C", "2007MNRAS.376..248H", "2008MNRAS.384..525M", "2009A&A...501..539U", "2009AJ....138..402S", "2009ApJS..185..289R", "2009PASP..121..591M", "2010AJ....139.1782L", "2010AJ....140.1868W", "2010ApJ...710..706M", "2010ApJS..189...37E", "2010MNRAS.405.1047G", "2010PASA...27..129F", "2011AJ....142...40M", "2011BSRSL..80..291W", "2012A&A...537A..10M", "2012AJ....143..149S", "2012AJ....144..166S", "2012IAUS..282..267S", "2014AJ....147..115F", "2015MNRAS.447.2322R" ]
[ "10.1093/mnras/stv1342", "10.48550/arXiv.1403.0975" ]
1403
1403.0975_arXiv.txt
In the more than 140 years since their first identification \citep{Wolf:jk}, Wolf--Rayet (WR) stars have remained one of the most interesting (and, at times, baffling) classes of stars. Their huge masses ($\ge25$~M\sun\ for initial stellar mass) and short lifetimes (typically $\sim3\times10^5$~years in the WR phase) make them excellent tracers of recent star formation, and their position in the stellar evolutionary chain is important to both stellar astrophysics and supernova theory. The intense stellar winds of WR stars add a significant amount of C and O, and some N, to the interstellar medium (ISM), and contribute a significant fraction of the ISM's energy and momentum budget; they also produce characteristic emission lines which give astronomers a probe into the atmospheres of these very hot, evolved stars. The most reliable method to date of WR detection has been to look for narrowband excess due to emission lines in the optical, particularly the strong \ion{He}{2} 4686~\AA\ line. However, probing the Milky Way using this technique is difficult beyond $\sim$3~kpc due to extreme dust extinction ($\sim$30~visual magnitudes across and through the centre of the Galaxy); in the near-infrared (NIR) only $\sim$3~magnitudes of extinction occurs across the Galactic plane \citep[see][section 7]{1999AJ....118..390S}. Clearly the NIR is the wavelength range of choice for searching out the vast majority of WR stars in the Milky Way. Over the last 5 years, advances in NIR selection techniques have led to the identification of more than 250 new WR stars, using a variety of methods. Two earlier papers in this series have focused on photometric techniques, and identified 112 new WR stars in the Galactic Plane \citep[][hereafter Paper I and Paper II, respectively]{2009AJ....138..402S,2012AJ....143..149S}. In addition, various colour-selection criteria in the near- and mid-infrared have been developed; see \citet{2007MNRAS.376..248H}, \citet{2009AAS...21460509M}, \citet{2011AJ....142...40M}, and \citet{2012A&A...537A..10M} (as well as \citet{2014AJ....147..115F} for a discussion of the specific techniques used in this paper). Simple models of the WR distribution in the Galaxy predict high concentrations near the Galactic centre (Paper I), where extremely crowded images cause usual photometric techniques to fail. However, using a new method of image subtraction, along with NIR and MIR colour-cuts we are able to significantly improve the selection criteria in the most crowded parts of the Galaxy; follow-up of candidates with this new technique produced 89 new emission-line sources in the Galactic Plane, including 49 new Wolf--Rayet stars. In section~\ref{sec:obs} we briefly describe the imaging survey and data reduction pipeline. Spectrographic follow-up data reductions are described in section~\ref{sec:spec}, and the new candidates are presented in section~\ref{sec:res}. The resulting overall distribution of all known WR stars is presented and discussed in section~\ref{sec:dist}. We briefly summarize our conclusions in section~\ref{sec:end}.
\label{sec:end} The imaging survey first introduced in Paper I has already produced 27 per cent of the known Galactic WR stars using photometric selection techniques \citep{2009AJ....138..402S,2012AJ....143..149S}. Using new reductions and image-subtraction methods, we have shown in this paper that there are still many more Galactic WR star candidates to be discovered and confirmed, particularly in the Southern hemisphere. The Galactic Centre region in particular has not been tapped, as the survey images are so crowded that special care must be taken with them. The pipeline described in this paper is capable of analysing all but the densest regions in our survey. The WR population simulations presented in Paper I predicted several thousand WR stars in the Milky Way, which have not yet been identified. However, the results presented here are not inconsistent with those predictions, for a few reasons. First, due to the necessity of spectroscopic follow-up to confirm each WR star, even with a high success rate the number of observations necessary is very large. Second, the majority of WR stars still unidentified lie either on the far side of the Galaxy, and are thus very faint, or near the Galactic Centre in prohibitively crowded fields. This publication makes use of data products from the Two Micron All Sky Survey, which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Centre/California Institute of Technology, funded by the National Aeronautics and Space Administration and the National Science Foundation. This publication 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. AFJM is grateful for financial assistance from NSERC (Canada) and FRQNT (Quebec). This research has made use of the NASA/ IPAC Infrared Science Archive, which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration. GK and MS gratefully acknowledge support from Ethel and Hilary Lipsitz. \nocite{*}
14
3
1403.0975
A new method of image subtraction is applied to images from a J, K, and narrow-band imaging survey of 300 deg<SUP>2</SUP> of the plane of the Galaxy, searching for new Wolf-Rayet (WR) stars. Our survey spans 150° in Galactic longitude and reaches b = ±1° with respect to the Galactic plane. The survey has a useful limiting magnitude of K = 15 over most of the observed Galactic plane, and K = 14 (due to severe crowding) within a few degrees of the Galactic Centre. The new image subtraction method described here (better than aperture or even point-spread-function photometry in very crowded fields) detected several thousand emission-line candidates. In 2011 and 2012 June and July, we spectroscopically followed up on 333 candidates with MDM-TIFKAM and Infrared Telescope Facility (IRTF)-SpeX, discovering 89 emission-line sources. These include 49 WR stars, 43 of them previously unidentified, including the most distant known Galactic WR stars, more than doubling the number on the far side of the Milky Way. We also demonstrate our survey's ability to detect very faint planetary nebulae and other NIR emission objects.
false
[ "Galactic WR", "Galactic longitude", "Galactic", "Infrared Telescope Facility", "K", "the observed Galactic plane", "severe crowding", "the Galactic plane", "the most distant known Galactic WR stars", "image subtraction", "several thousand emission-line candidates", "the Galactic Centre", "very faint planetary nebulae and other NIR emission objects", "89 emission-line sources", "49 WR stars", "respect", "Galaxy", "b", "the Milky Way", "The new image subtraction method" ]
9.592138
10.230614
150
810315
[ "Zhang, L. Y.", "He, J. J.", "Parikh, A.", "Xu, S. W.", "Yamaguchi, H.", "Kahl, D.", "Kubono, S.", "Mohr, P.", "Hu, J.", "Ma, P.", "Chen, S. Z.", "Wakabayashi, Y.", "Wang, H. W.", "Tian, W. D.", "Chen, R. F.", "Guo, B.", "Hashimoto, T.", "Togano, Y.", "Hayakawa, S.", "Teranishi, T.", "Iwasa, N.", "Yamada, T.", "Komatsubara, T.", "Zhang, Y. H.", "Zhou, X. H." ]
2014PhRvC..89a5804Z
[ "Investigation of the thermonuclear <SUP>18</SUP>Ne(α,p)<SUP>21</SUP>Na reaction rate via resonant elastic scattering of <SUP>21</SUP>Na + p" ]
19
[ "Chinese Academy of Sciences, Lanzhou 730000, China, Institute of Modern Physics, Chinese Academy of Sciences, Lanzhou 730000, China", "Chinese Academy of Sciences, Lanzhou 730000, China, Institute of Modern Physics, Chinese Academy of Sciences, Lanzhou 730000, China", "EUETIB, Universitat Politècnica de Catalunya, Barcelona E-08036, Spain, Departament de Física i Enginyeria Nuclear, EUETIB, Universitat Politècnica de Catalunya, Barcelona E-08036, Spain; , Barcelona E-08034, Spain, Institut d'Estudis Espacials de Catalunya, Barcelona E-08034, Spain", "Chinese Academy of Sciences, Lanzhou 730000, China, Institute of Modern Physics, Chinese Academy of Sciences, Lanzhou 730000, China; , Beijing 100049, China, University of Chinese Academy of Sciences, Beijing 100049, China", "Department of Physics, University of Tsukuba, Ibaraki 305-8571, Japan", "Department of Physics, University of Tsukuba, Ibaraki 305-8571, Japan", "Chinese Academy of Sciences, Lanzhou 730000, China, Institute of Modern Physics, Chinese Academy of Sciences, Lanzhou 730000, China; Department of Physics, University of Tsukuba, Ibaraki 305-8571, Japan; , Wako, Saitama 351-0198, Japan, RIKEN (Institute of Physical and Chemical Research), Wako, Saitama 351-0198, Japan", "Schwäbisch Hall D-74523, Germany, Diakonie-Klinikum, Schwäbisch Hall D-74523, Germany; , Debrecen H-4001, Hungary, Institute of Nuclear Research (ATOMKI), Debrecen H-4001, Hungary", "Chinese Academy of Sciences, Lanzhou 730000, China, Institute of Modern Physics, Chinese Academy of Sciences, Lanzhou 730000, China", "Chinese Academy of Sciences, Lanzhou 730000, China, Institute of Modern Physics, Chinese Academy of Sciences, Lanzhou 730000, China", "Chinese Academy of Sciences, Lanzhou 730000, China, Institute of Modern Physics, Chinese Academy of Sciences, Lanzhou 730000, China; , Beijing 100049, China, University of Chinese Academy of Sciences, Beijing 100049, China", "Japan Atomic Energy Agency, Ibaraki 319-1106, Japan, Advanced Science Research Center, Japan Atomic Energy Agency, Ibaraki 319-1106, Japan", "Chinese Academy of Sciences, Shanghai 201800, China, Shanghai Institute of Applied Physics, Chinese Academy of Sciences, Shanghai 201800, China", "Chinese Academy of Sciences, Shanghai 201800, China, Shanghai Institute of Applied Physics, Chinese Academy of Sciences, Shanghai 201800, China", "Chinese Academy of Sciences, Lanzhou 730000, China, Institute of Modern Physics, Chinese Academy of Sciences, Lanzhou 730000, China", "Post Office Box 275(46), Beijing 102413, China, China Institute of Atomic Energy, Post Office Box 275(46), Beijing 102413, China", "Department of Physics, University of Tsukuba, Ibaraki 305-8571, Japan", "Wako, Saitama 351-0198, Japan, RIKEN (Institute of Physical and Chemical Research), Wako, Saitama 351-0198, Japan", "Department of Physics, University of Tsukuba, Ibaraki 305-8571, Japan", "Department of Physics, University of Tsukuba, Ibaraki 305-8571, Japan", "Department of Physics, University of Tsukuba, Ibaraki 305-8571, Japan", "Department of Physics, University of Tsukuba, Ibaraki 305-8571, Japan", "Department of Physics, University of Tsukuba, Ibaraki 305-8571, Japan", "Chinese Academy of Sciences, Lanzhou 730000, China, Institute of Modern Physics, Chinese Academy of Sciences, Lanzhou 730000, China", "Chinese Academy of Sciences, Lanzhou 730000, China, Institute of Modern Physics, Chinese Academy of Sciences, Lanzhou 730000, China" ]
[ "2014PhRvC..90f5806M", "2015NDS...127...69B", "2016EPJWC.11709005Y", "2018AIPC.1947b0022Y", "2018JKPS...73.1055C", "2018JPhG...45i3001M", "2018PhRvC..97a5802K", "2020JPhCS1643a2069Y", "2020NPNew..30...21Y", "2021APPSB..31...18A", "2021PhRvC.103e5815D", "2021PhRvL.127q2701H", "2021PrPNP.11903857S", "2021arXiv210906537B", "2022JPhG...49k0502S", "2022PhRvC.105e5806A", "2023EPJA...59..112K", "2023FrP....1163299C", "2023PhRvC.107e5802B" ]
[ "astronomy", "physics" ]
17
[ "25.60.-t", "23.50.+z", "26.50.+x", "27.30.+t", "Reactions induced by unstable nuclei", "Decay by proton emission", "Nuclear physics aspects of novae supernovae and other explosive environments", "20&lt", "=A&lt", "=38", "Astrophysics - Solar and Stellar Astrophysics", "Nuclear Experiment" ]
[ "1958RvMP...30..257L", "1981ApJS...45..389W", "1985NIMPA.236..128N", "1990NuPhA.521....1E", "1991PhRvC..44.2776G", "1992ARNPS..42...39C", "1993SSRv...62..223L", "1995PhRvC..51..392G", "1998PhR...294..167S", "1999JPhG...25R.133W", "1999PhRvC..59.3402B", "2000ADNDT..75....1R", "2001NIMPA.470..562K", "2001NuPhA.693..221K", "2001PhRvC..63f5807C", "2002EPJA...13..217K", "2002PhRvC..66a5804C", "2002PhRvC..66e5802G", "2003NuPhA.718..608B", "2003PhLB..556...27T", "2004ApJS..151...75W", "2004PhRvL..93o0801M", "2005NIMPA.539...74Y", "2005PhRvC..71b5802R", "2007PhLB..650..129T", "2007PhRvC..76e5802H", "2007nps..book.....I", "2008ApJS..178..110P", "2008EPJA...36....1H", "2008NIMPA.589..150Y", "2009PhLB..672..230Y", "2009PhRvC..79d5802P", "2009PhRvC..79e5804C", "2009PhRvC..80a5801H", "2009PhRvC..80e5804M", "2010NuPhA.841....1L", "2012ChPhC..36....3M", "2012PhRvL.108x2701S", "2013PhRvC..87c5801M", "2013PhRvC..88a2801H", "2013PrPNP..69..225P" ]
[ "10.1103/PhysRevC.89.015804", "10.48550/arXiv.1403.4668" ]
1403
1403.4668.txt
\label{sec:introduction} Explosive hydrogen and helium burning are thought to be the main source of energy generation and nuclear trajectory to higher mass on the proton-rich side of the nuclear chart in type I x-ray bursts (XRBs)~\cite{champage,wiescher,bib:lew93,bib:str06,bib:par13}. XRBs are characterized by a sudden increase of x-ray emission within only a few seconds to a total energy output of about 10$^{40}$ ergs, which is observed to repeat with some regularity. The recurrence time for single bursts can range from hours to days at the typical temperature of 0.4--2 GK. The bursts have been interpreted as being generated by thermonuclear runaway on the surface of a neutron star that accretes H- and He-rich material from a less evolved companion star in a close binary system. In XRBs, the hydrogen burning initially occurs via the hot CNO cycle (HCNO): \begin{center} \begin{spacing}{2.0} $^{12}$C($p$,$\gamma$)$^{13}$N($p$,$\gamma$)$^{14}$O($e^{+}\nu$)$^{14}$N($p$,$\gamma$)$^{15}$O($e^{+}\nu$) $^{15}$N($p$,$\alpha$)$^{12}$C, \end{spacing} \end{center} while the $^{13}$N($e^{+}\nu$)$^{13}$C $\beta$-decay in the CNO cycle is bypassed by the $^{13}$N($p$,$\gamma$)$^{14}$O reaction. The temperature of the accretion envelope increases as the compressing and exothermic nuclear reactions going on. When the temperature reaches about 0.4 GK, the second HCNO cycle becomes dominant: \begin{center} \begin{spacing}{2.0} $^{12}$C($p$,$\gamma$)$^{13}$N($p$,$\gamma$)$^{14}$O($\alpha$,$p$)$^{17}$F($p$,$\gamma$)$^{18}$Ne($e^{+}\nu$)$^{18}$F($p$,$\alpha$) $^{15}$O($e^{+}\nu$)$^{15}$N($p$,$\alpha$)$^{12}$C. \end{spacing} \end{center} It was predicted~\cite{champage,wiescher} that the $^{18}$Ne waiting point nucleus in the second HCNO cycle could be bypassed by the $^{18}$Ne($\alpha$,$p$)$^{21}$Na reaction at $T$$\approx$0.6 GK, and subsequently, the reaction chain breaks out, eventually leading to the rp-process~\cite{wal81,sch98,woo04}. However, over stellar temperatures achieved in XRBs, this rate has not been sufficiently well determined. The thermonuclear $^{18}$Ne($\alpha$,$p$)$^{21}$Na rate is thought to be dominated by contributions from resonances in the compound nucleus $^{22}$Mg above the $\alpha$ threshold at $Q_\alpha$=8.142 MeV~\cite{wang}. As for XRBs, the temperature region of interest is about 0.4--2.0 GK, corresponding to an excitation region of $E_x$=8.6--11.0 MeV in $^{22}$Mg. G\"{o}rres {\it et al.} made the first estimate~\cite{gorres} of this rate with rather limited experimental level-structure information in $^{22}$Mg. The energies for the $^{22}$Mg resonances were estimated simply by shifting those of known natural-parity states in the mirror $^{22}$Ne by a fixed amount (about 200 keV). The uncertainty of this first rate was mainly caused by the errors in resonant energies (or excitation energies) and resonant strengths of the excited states above the $\alpha$ threshold in $^{22}$Mg. After that, the precise locations of the excited states in $^{22}$Mg were studied extensively by many transfer reaction experiments. For example, Chen {\it et al.}~\cite{chen} determined the excitation energies with a typical uncertainty of 20--30 keV in a $^{12}$C($^{16}$O,$^{6}$He)$^{22}$Mg experiment. However, the spin-parity assignments assumed and the spectroscopic $S_\alpha$ factors adopted following the idea of G\"{o}rres {\it et al.} were still uncertain. Caggiano {\it et al.}~\cite{caggiano} and Berg {\it et al.}~\cite{berg} measured the excitation energies with a better precision (about 10--20 keV), but no spin-parity assignment was given. Later, Matic {\it et al.}~\cite{matic} measured the excitation energies precisely by a $^{24}$Mg($p$,$t$)$^{22}$Mg experiment, with uncertainty of about 1--15 keV achieved for most states above the $\alpha$ threshold; the spin-parity assignments were tentatively made based on the shell-model calculation or those of mirror states in $^{22}$Ne. Thus, the thermonuclear $^{18}$Ne($\alpha$,$p$)$^{21}$Na rate was constrained very well in the resonant energy aspect. In a later $^{24}$Mg($p$,$t$)$^{22}$Mg measurement, Chae {\it et al.}~\cite{chae} observed six excited states in $^{22}$Mg above the $\alpha$ threshold, and some spin-parity assignments were made via an angular distribution measurement. However, the insufficient resolution of their measurement at the center-of-mass (c.m.) scattering angles of $\theta_{c.m.}$ above 20${^\circ}$ made such $J^{\pi}$ assignments questionable~\cite{matic} ({\it e.g.}, the 8.495 MeV peak was contaminated by the nearby 8.572 and 8.658 MeV states as shown in their Fig. 3). In our previous experiment of $^{21}$Na+$p$ resonant elastic scattering~\cite{he_epja,he}, new spin-parity assignments were made only tentatively for the the 8.547 and 8.614 MeV states in $^{22}$Mg due to low statistics. Those assignments gave a quite different rate for the $^{18}$Ne($\alpha$,$p$)$^{21}$Na reaction compared to the rate estimated in Chen {\it et al.} work~\cite{chen}. Such tentative assignments clearly motivate further investigation. A comparison of all available reaction rates of $^{18}$Ne($\alpha$,$p$)$^{21}$Na shows discrepancies of up to several orders of magnitude around $T$$\sim$1 GK~\cite{matic}. So far, more than 40 levels (up to $E_x$=13.01 MeV) have been observed above the $\alpha$ threshold in $^{22}$Mg. Such high level density suggests that a statistical model approach might provide a reliable estimate of the rate. However, only natural-parity states in $^{22}$Mg can be populated by the $^{18}$Ne+$\alpha$ channel, and thus the effective level density will be considerably lower. It remains unclear whether the statistical-model calculations provide a reliable rate estimation in a wide temperature region~\cite{matic}. There are still many resonances (above the $\alpha$ threshold) without firm spin-parity assignments, which need to be determined experimentally. As a consequence, the accuracy of the current $^{18}$Ne($\alpha$,$p$)$^{21}$Na reaction rate is mainly limited by the lack of experimental spin-parities and $\alpha$ partial widths $\Gamma_\alpha$ (or spectroscopic factors $S_\alpha$) of the resonances in $^{22}$Mg above the $\alpha$ threshold. So far, only two direct measurements~\cite{bradfield,groombridge} for the $^{18}$Ne($\alpha$,$p$)$^{21}$Na reaction were reported. The lowest energies achieved in these studies ($E_{c.m.}$=2.0 and 1.7 MeV) are still too high compared with the energy region $E_{c.m.}$$\leq $1.5 MeV of interest for HCNO breakout in XRBs. New results~\cite{salter} have recently become available, which determined the $^{18}$Ne($\alpha$,$p_0$)$^{21}$Na cross sections in the energy region of $E_{c.m.}$=1.19--2.57 MeV by measuring those of the time-reversal reaction $^{21}$Na($p$,$\alpha$)$^{18}$Ne in inverse kinematics. In addition, similar experiments were performed at the Argonne National Laboratory (ANL), but the results were only reported in the ANL annual reports~\cite{anl}. The ANL cross section data are consistent with those in Ref.~\cite{salter}. Nonetheless, these results are still insufficient for a reliable rate calculation at all temperatures encountered within XRBs. Recently, a new reaction rate was recommended based on the combined analysis of all literature data~\cite{matic2}. In this work, the $^{18}$Ne($\alpha$,$p$)$^{21}$Na rate is determined via the measurement of the resonant elastic scattering of $^{21}$Na+$p$. This is an entirely new high-statistics experiment compared to the previous one~\cite{he_epja,he}. In the resonant elastic-scattering mechanism, $^{22}$Mg is formed via the fusion of $^{21}$Na+$p$ as an excited compound nucleus, whose states promptly decay back into $^{21}$Na+$p$. This process interferes with Coulomb scattering resulting in a characteristic resonance pattern in the excitation function~\cite{ruiz}. With this approach, the excitation function was obtained simultaneously in a wide range of 5.5--9.2 MeV in $^{22}$Mg with a well-established thick-target method~\cite{art90,gal91,kub01}. In total, 23 states above the proton-threshold in $^{22}$Mg were observed, and their resonant parameters were determined via an $R$-matrix analysis of the experimental data. Part of the experimental results previously reported in Ref.~\cite{he_prc_rc} is revisited through a more detailed analysis. The detailed experimental results presented here supersede those of Ref.~\cite{he_prc_rc}.
We have studied the resonant elastic-scattering of $^{21}$Na+$p$ using a radioactive ion beam of $^{21}$Na with the thick-target method. The $E_{c.m.}$ spectra were reconstructed from the inverse kinematics. In total, 23 resonances above the proton-threshold in $^{22}$Mg were observed. The relevant proton resonant parameters have been determined by the $R$-matrix analysis of the center-of-mass differential cross-section data at different scattering angles. The $^{18}$Ne($\alpha$,$p$)$^{21}$Na reaction rate is recalculated with the present parameters. A new recommended rate is given by combining the results from different experimental techniques, and our new rate deviates considerably from the recent recommended rate of Mohr \& Matic below $\sim$0.55 GK. The astrophysical impact of our new rate has been investigated through one-zone postprocessing x-ray burst calculations. Compared to previous rates in Refs.~\cite{gorres,chen,chae}, the new rate increases the energy production rate by factors of 1.4--1.8 at early time (between 0.3--0.4 s, or equivalently, between 0.6--0.9 GK during the burst) of the burst, and the $^{18}$Ne($\alpha$,$p$)$^{21}$Na reaction flux is also enhanced about two times at that time. The breakout onset temperature for this reaction occurs at around 0.57 GK, lowered by 0.11 GK due to the increase of the reaction rate. Despite the different $J^\pi$ values adopted in the present and Matic {\it et al.} $^{18}$Ne($\alpha$,$p$)$^{21}$Na rate calculations (and the consequent differences in deduced thermonuclear rates), our models give very similar XRB nuclear energy generation rates. This suggests that $J^\pi$ values for relevant states in $^{22}$Mg are, for the moment, sufficiently well known for our models. Future measurements should primarily focus on measuring other quantities of interest (such as spectroscopic factors, partial widths or the precise cross section data), which can further constrain this rate. \begin{center} \textbf
14
3
1403.4668
The <SUP>18</SUP>Ne(α,p)<SUP>21</SUP>Na reaction is thought to be one of the key breakout reactions from the hot CNO cycles to the r<SUB>p process in type I x-ray bursts. In this work, the resonant properties of the compound nucleus <SUP>22</SUP>Mg have been investigated by measuring the resonant elastic scattering of <SUP>21</SUP>Na + p. An 89-MeV <SUP>21</SUP>Na radioactive beam delivered from the CNS Radioactive Ion Beam Separator bombarded an 8.8 mg/cm<SUP>2</SUP> thick polyethylene (CH2</SUB>)n target. The <SUP>21</SUP>Na beam intensity was about 2×10<SUP>5</SUP> pps, with a purity of about 70% on target. The recoiled protons were measured at the center-of-mass scattering angles of θc<SUB>.m.</SUB>≈175.2<SUP>∘</SUP>, 152.2<SUP>∘</SUP>, and 150.5<SUP>∘</SUP> by three sets of ΔE-E telescopes, respectively. The excitation function was obtained with the thick-target method over energies E<SUB>x</SUB>(<SUP>22</SUP>Mg)=5.5-9.2 MeV. In total, 23 states above the proton-threshold in <SUP>22</SUP>Mg were observed, and their resonant parameters were determined via an R-matrix analysis of the excitation functions. We have made several new J<SUP>π</SUP> assignments and confirmed some tentative assignments made in previous work. The thermonuclear <SUP>18</SUP>Ne(α,p)<SUP>21</SUP>Na rate has been recalculated based on our recommended spin-parity assignments. The astrophysical impact of our new rate has been investigated through one-zone postprocessing x-ray burst calculations. We find that the <SUP>18</SUP>Ne(α,p)<SUP>21</SUP>Na rate significantly affects the peak nuclear energy generation rate, reaction fluxes, and onset temperature of this breakout reaction in these astrophysical phenomena.
false
[ "type I x-ray bursts", "CH2</SUB>)n target", "target", "> assignments", "E", "previous work", "reaction fluxes", "energies", "-", "one-zone postprocessing x-ray burst calculations", "the peak nuclear energy generation rate", "<", "Mg)=5.5", "SUP>∘</SUP", "ΔE-E telescopes", "several new J<SUP>π</SUP", "some tentative assignments", "the key breakout reactions", "the thick-target method", "this breakout reaction" ]
3.522935
2.426963
57
438023
[ "Amorín, R.", "Sommariva, V.", "Castellano, M.", "Grazian, A.", "Tasca, L. A. M.", "Fontana, A.", "Pentericci, L.", "Cassata, P.", "Garilli, B.", "Le Brun, V.", "Le Fèvre, O.", "Maccagni, D.", "Thomas, R.", "Vanzella, E.", "Zamorani, G.", "Zucca, E.", "Bardelli, S.", "Capak, P.", "Cassará, L. P.", "Cimatti, A.", "Cuby, J. G.", "Cucciati, O.", "de la Torre, S.", "Durkalec, A.", "Giavalisco, M.", "Hathi, N. P.", "Ilbert, O.", "Lemaux, B. C.", "Moreau, C.", "Paltani, S.", "Ribeiro, B.", "Salvato, M.", "Schaerer, D.", "Scodeggio, M.", "Talia, M.", "Taniguchi, Y.", "Tresse, L.", "Vergani, D.", "Wang, P. W.", "Charlot, S.", "Contini, T.", "Fotopoulou, S.", "López-Sanjuan, C.", "Mellier, Y.", "Scoville, N." ]
2014A&A...568L...8A
[ "Discovering extremely compact and metal-poor, star-forming dwarf galaxies out to z ~ 0.9 in the VIMOS Ultra-Deep Survey" ]
53
[ "INAF-Osservatorio Astronomico di Roma, via di Frascati 33, 00040, Monte Porzio Catone, Italy", "University of Bologna, Department of Physics and Astronomy (DIFA), V.le 6/2, 40127, Berti Pichat, Bologna; INAF-Osservatorio Astronomico di Roma, via di Frascati 33, 00040, Monte Porzio Catone, Italy", "INAF-Osservatorio Astronomico di Roma, via di Frascati 33, 00040, Monte Porzio Catone, Italy", "INAF-Osservatorio Astronomico di Roma, via di Frascati 33, 00040, Monte Porzio Catone, Italy", "Aix Marseille Université, CNRS, LAM (Laboratoire d'Astrophysique de Marseille) UMR 7326, 13388, Marseille, France", "INAF-Osservatorio Astronomico di Roma, via di Frascati 33, 00040, Monte Porzio Catone, Italy", "INAF-Osservatorio Astronomico di Roma, via di Frascati 33, 00040, Monte Porzio Catone, Italy", "Aix Marseille Université, CNRS, LAM (Laboratoire d'Astrophysique de Marseille) UMR 7326, 13388, Marseille, France", "INAF-IASF, via Bassini 15, 20133, Milano, Italy", "Aix Marseille Université, CNRS, LAM (Laboratoire d'Astrophysique de Marseille) UMR 7326, 13388, Marseille, France", "Aix Marseille Université, CNRS, LAM (Laboratoire d'Astrophysique de Marseille) UMR 7326, 13388, Marseille, France", "INAF-IASF, via Bassini 15, 20133, Milano, Italy", "Aix Marseille Université, CNRS, LAM (Laboratoire d'Astrophysique de Marseille) UMR 7326, 13388, Marseille, France", "INAF-Osservatorio Astronomico di Bologna, via Ranzani 1, 40127, Bologna, Italy", "INAF-Osservatorio Astronomico di Bologna, via Ranzani 1, 40127, Bologna, Italy", "INAF-Osservatorio Astronomico di Bologna, via Ranzani 1, 40127, Bologna, Italy", "INAF-Osservatorio Astronomico di Bologna, via Ranzani 1, 40127, Bologna, Italy", "Department of Astronomy, California Institute of Technology, 1200 E. California Blvd., MC 249-17, Pasadena, CA, 91125, USA", "INAF-IASF, via Bassini 15, 20133, Milano, Italy", "University of Bologna, Department of Physics and Astronomy (DIFA), V.le 6/2, 40127, Berti Pichat, Bologna", "Aix Marseille Université, CNRS, LAM (Laboratoire d'Astrophysique de Marseille) UMR 7326, 13388, Marseille, France", "University of Bologna, Department of Physics and Astronomy (DIFA), V.le 6/2, 40127, Berti Pichat, Bologna; INAF-Osservatorio Astronomico di Bologna, via Ranzani 1, 40127, Bologna, Italy", "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", "Astronomy Department, University of Massachusetts, Amherst, MA, 01003, USA", "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", "Aix Marseille Université, CNRS, LAM (Laboratoire d'Astrophysique de Marseille) UMR 7326, 13388, Marseille, France", "Department of Astronomy, University of Geneva, ch. d'cogia 16, 1290, Versoix", "Aix Marseille Université, CNRS, LAM (Laboratoire d'Astrophysique de Marseille) UMR 7326, 13388, Marseille, France", "Max-Planck-Institut für Extraterrestrische Physik, Postfach 1312, 85741, Garching bei München, Germany", "Geneva Observatory, University of Geneva, ch. des Maillettes 51, CH-1290, Versoix, Switzerland; Institut de Recherche en Astrophysique et Planétologie - IRAP, CNRS, Université de Toulouse, UPS-OMP, 14 avenue E. Belin, 31400, Toulouse, France", "INAF-IASF, via Bassini 15, 20133, Milano, Italy", "University of Bologna, Department of Physics and Astronomy (DIFA), V.le 6/2, 40127, Berti Pichat, Bologna", "Research Center for Space and Cosmic Evolution, Ehime University, Bunkyo-cho 2-5, 790-8577, Matsuyama, Japan", "Aix Marseille Université, CNRS, LAM (Laboratoire d'Astrophysique de Marseille) UMR 7326, 13388, Marseille, France", "INAF-IASF Bologna, via Gobetti 101, 40129, Bologna, Italy; INAF-Osservatorio Astronomico di Bologna, via Ranzani 1, 40127, Bologna, Italy", "Aix Marseille Université, CNRS, LAM (Laboratoire d'Astrophysique de Marseille) UMR 7326, 13388, Marseille, France", "Institut d'Astrophysique de Paris, UMR7095 CNRS, Université Pierre et Marie Curie, 98 bis Boulevard Arago, 75014, Paris, France", "Institut de Recherche en Astrophysique et Planétologie - IRAP, CNRS, Université de Toulouse, UPS-OMP, 14 avenue E. Belin, 31400, Toulouse, France", "Department of Astronomy, University of Geneva, ch. d'cogia 16, 1290, Versoix", "Centro de Estudios de Física del Cosmos de Aragón, Teruel, Spain", "Institut d'Astrophysique de Paris, UMR7095 CNRS, Université Pierre et Marie Curie, 98 bis Boulevard Arago, 75014, Paris, France", "Department of Astronomy, California Institute of Technology, 1200 E. California Blvd., MC 249-17, Pasadena, CA, 91125, USA" ]
[ "2014A&A...572A..41L", "2014ApJ...788L...4A", "2014ApJ...791...17M", "2014MNRAS.442.1794O", "2014Msngr.155...37L", "2014PhDT........11S", "2015A&A...576A..79L", "2015A&A...578A.105A", "2015A&A...583A..88N", "2015ApJ...799...36R", "2015ApJ...805...45L", "2015ApJ...813..126J", "2015MNRAS.449..917L", "2015MNRAS.452.1567C", "2015Natur.521...54Z", "2015PASJ...67...80H", "2015sf2a.conf...77D", "2016A&A...585A..51D", "2016A&A...588A..23A", "2016A&A...592A.122H", "2016ApJ...819...81R", "2016ApJ...819..110S", "2016ApJ...822..103G", "2016ApJ...828...67L", "2017A&A...599A..28K", "2017A&A...601A..95C", "2017ASSL..430...67S", "2017RAA....17...41G", "2018ApJ...854...29M", "2018ApJ...869...15G", "2018MNRAS.473.1258S", "2018MNRAS.476.4765S", "2018PASJ...70S..35M", "2019ApJ...871..197D", "2019ApJ...872..145J", "2019ApJ...882...17Y", "2019ApJ...882..116C", "2020A&A...633A..68D", "2020MNRAS.491.2254W", "2020MNRAS.494..719M", "2020MNRAS.494.3541H", "2020arXiv200511883E", "2021ApJ...910..135S", "2021MNRAS.503.4105T", "2021MNRAS.504.1237P", "2022ApJ...935..150L", "2022MNRAS.514.2627C", "2023A&A...676A..53L", "2023ApJ...952...76M", "2023MNRAS.522.2089D", "2023arXiv231207799D", "2024ApJ...961..173Z", "2024arXiv240305362L" ]
[ "astronomy" ]
5
[ "galaxies: evolution", "galaxies: high-redshift", "galaxies: starburst", "galaxies: abundances", "galaxies: dwarf", "galaxies: fundamental parameters", "Astrophysics - Astrophysics of Galaxies", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
[ "1991ApJ...380..140M", "1996AJ....112..839C", "1998ApJ...498..541K", "2000ApJ...533..682C", "2003ApJ...596..797F", "2003MNRAS.344.1000B", "2003MNRAS.346..105P", "2003PASP..115..763C", "2003SPIE.4841.1670L", "2004ApJS..153..429K", "2005ApJ...635L..21H", "2006A&A...448..907L", "2007AJ....133..882K", "2008A&A...488..463M", "2008A&A...491..113P", "2008ApJ...675..234P", "2008MNRAS.383..209H", "2009A&A...495...53L", "2009A&A...503..379T", "2009A&A...505...63G", "2009ApJ...698.2014H", "2009MNRAS.398..949P", "2009MNRAS.399.1191C", "2010ApJ...715L.128A", "2010MNRAS.408.2115M", "2011ApJ...728...38Y", "2011ApJ...730..137Z", "2011ApJ...736..104J", "2011ApJ...742..111V", "2011ApJ...743...77M", "2011ApJ...743..121A", "2011MNRAS.414.1263M", "2012ApJ...749..185A", "2012ApJ...754L..29W", "2013A&A...549A..25P", "2013ApJ...763..145D", "2013ApJ...765..140A", "2013ApJ...766...91J", "2013ApJ...769..148H", "2014A&A...563A..58T", "2014A&A...566A..19C", "2014ApJ...780..122L", "2014ApJ...783...45S", "2014ApJ...788L...4A", "2014ApJ...791...17M", "2014MNRAS.442..900N", "2015A&A...576A..79L", "2015A&A...578A.105A" ]
[ "10.1051/0004-6361/201423816", "10.48550/arXiv.1403.3692" ]
1403
1403.3692_arXiv.txt
\label{sect:intro} Over the last 8 billion years % a large fraction of low-mass (M$_{\star}\la$\,10$^{9}$\,M$_{\odot}$) galaxies are still seen rapidly assembling most of their present-day stellar mass \citep{Cowie1996,Perez-Gonzalez2008}. Tracing the spectrophotometric properties of these vigorous star-forming \textit{dwarf} galaxies (SFDGs) out to $z\sim$1 is essential not only to study how they evolve through cosmic time, but also to understand the physical mechanisms driving the first stages of stellar mass build-up and chemical enrichment. To this end, key insights % can be obtained from the tight relations found between stellar mass, metallicity and star formation rate (SFR). However, the shape and normalization of these relations at different redshifts are still poorly constrained at their low-mass end. While in the local Universe the mass-metallicity relation (MZR) has been extended down to 10$^8$\,M$_{\odot}$ \citep[e.g.][]{Andrews2013}, at intermediate and high redshifts, dwarf galaxies are strongly underrepresented \citep[e.g.][]{Henry2013}. These SFDGs are usually identified by their blue colors, high surface brightness and strong emission-lines. They include a rare population of extreme emission-line galaxies (EELGs) with the largest nebular content and lowest metal abundances \citep[e.g.][]{Kniazev2004,Papaderos2008,Hu2009,Atek2011,Morales-Luis2011}. Due to their high equivalent widths (EWs), an increasing number of EELGs are being discovered and characterized by deep spectroscopic surveys out to $z\sim$1 \citep[e.g.][]{Hoyos2005,Ly2014,Amorin2014a} and beyond \citep[e.g.][]{vdWel2011,Maseda2014}. In this \textit{Letter} we report the discovery of a sample of 31 EELGs at $0.2\la z \la 0.9$ identified from the \textit{VIMOS Ultra-Deep Survey} \citep[VUDS;][]{LeFevre2014}. We study their physical properties as part of a larger, ongoing study aimed at investigating the evolution of SFDGs out to $z\sim$\,1 using very deep spectroscopy \citep[e.g.][]{Amorin2014a}. The sensitivity of our VUDS spectra, detecting emission lines as faint as $\sim$1.5$\times$\,10$^{-18}$\,erg s$^{-1}$ cm$^{-2}$, % makes it possible e.g., to derive $T_e$-based metallicities for a fraction of such faint galaxies. Thus, the present sample extends previous studies of star-forming (SF) galaxies at similar redshifts in size and limiting magnitude \citep{Henry2013,Ly2014}, allowing us to study in greater detail the LZR and MZR at $z<1$ two decades below 10$^{9}$M$_{\odot}$ with galaxies showing a wide range of properties, including a number of extremely metal-poor galaxies. % Throughout this paper we adopt a standard $\Lambda$-CDM cosmology with $h$ = 0.7, $\Omega_m$ = 0.3 and $\Omega_\Lambda$ = 0.7.
In Fig.~\ref{M-SFR} we show the SFR-mass diagram for the EELGs in VUDS and from the literature. Star formation rates are derived from the extinction-corrected \ha\ or \hb\ luminosities using the calibration of \citet{Kennicutt1998} and assuming a \citet{Chabrier2003} IMF. At a given redshift, our EELGs show SFRs and stellar masses a factor of $\sim$10 lower than similar samples from the literature. However, nearly all EELGs shown in Fig.~\ref{M-SFR} are well above the extrapolation to low stellar mass of the main sequence of galaxies \citep{Whitaker2012} at a given $z$. The EELGs in VUDS show enhanced \textit{specific} SFRs (sSFR$\sim$10$^{-9}$-10$^{-7}$\,yr$^{-1}$) and SFR surface densities (median $\Sigma_{\rm SFR}=$ SFR$/2\pi r^2_{50} =$\,0.35 ($\sigma =$\,0.19)\,M$_{\odot}$\,yr$^{-1}$\,kpc$^{-2}$), comparable to more luminous galaxy-wide starbursts at similar and higher redshifts \citep{Ly2014,Amorin2014a,Amorin2014b}. \begin{figure}[t!] \centering \includegraphics[angle=0,width=7.6cm]{Fig5a_rev.eps}\\\vspace{2mm} \includegraphics[angle=0,width=7.6cm]{Fig5b_rev.eps}\\\vspace{2mm} \includegraphics[angle=0,width=7.6cm]{Fig5c_rev.eps}\\\vspace{2mm} \includegraphics[angle=0,width=7.6cm]{Fig5d_rev.eps}\\ \caption{{($a$) Luminosity-metallicity and ($b$) mass-metallicity relations for VUDS EELGs and SFDGs from the literature. Metallicity differences with respect to the extrapolation to low stellar mass of the FMR by \citet[][]{Andrews2013} and \citet[][]{Mannucci2011} are shown in $(c)$ and $(d)$, respectively. Dashed lines indicate 1$\sigma$ deviations for these relations. % Colors and symbols % are as in Fig.~\ref{M-SFR}. The data have been homogenized to the Chabrier IMF and the same strong-line metallicity calibration presented in Section~3.1.} } \label{MZR} \end{figure} In Fig.~\ref{MZR} we study the LZR and MZR traced by EELGs in VUDS and other low-mass galaxies at $0<z<1$. The EELGs {extend} the LZR down to $M_{\rm AB}(B) \sim -14.5$ and the MZR down to M$_{\star} \sim$\,10$^{7}$M$_{\odot}$, which means $\sim$1 dex lower than previous studies \citep[e.g.,][]{Henry2013}, thus increasing substantially the number of low-mass galaxies under study, especially at $z$\,$\ga$\,0.5. Despite the relatively large scatter, VUDS EELGs appear to follow the LZR and MZR of more luminous and massive SFDGs. In particular, we find most EELGs in broad agreement with the local ($z<0.3$) LZR of \citet{Guseva2009} and MZR of \citet{Andrews2013}, which have been derived from galaxies with {$T_e$-based metallicities. There is nevertheless a tendency for EELGs with larger EWs % to be more metal-poor at a given luminosity, stellar mass, and redshift.} These galaxies are those with the highest sSFR, i.e., those with the largest deviations from the main sequence of star formation at a given $z$, shown in Fig.~\ref{M-SFR}. While they follow more reliably the LZR traced by extremely metal-poor galaxies \citep[e.g.,][]{Kewley2007,Hu2009}, they tend to lie below the local MZR, similarly to other extreme galaxies \citep[see, e.g., the \textit{green peas}][]{Amorin2010}. Part of the above apparent dependence of the MZR on SFR can be explained in terms of the fundamental metallicity relation \citep[FMR;][]{Mannucci2010}, which suggests that galaxies with higher sSFR tend to be more metal-poor at a given stellar mass. As shown in Fig.~\ref{MZR}, the position of the VUDS EELGs appears broadly consistent with the extrapolation to low masses of the FMR, independently of the parametrization and metallicity scale adopted. We notice, however, that the scatter in the FMR for EELGs ($\sigma \sim$\,0.20) seems slightly larger than expected for magnitude-selected samples in the local universe. Overall, the above results are consistent with a picture where the most extreme SFDGs are very gas-rich galaxies experiencing an early stage of a galaxy-wide starburst, possibly fed by recent accretion of metal-poor gas \citep[e.g.,][]{Amorin2010,SanchezAlmeida2014}. In this picture, at least part of the scatter in the above scaling relations could be produced by differences in the accretion and star formation histories. Figure~\ref{MZR} also suggests that the shape of the MZR can be very sensitive to selection effects in its very low-mass end. Gas-rich dwarfs with prominent emission lines, enhanced sSFR and low metallicities may be overrepresented with respect to the global population of SFDGs in magnitude-selected spectroscopic samples at these redshifts, making the shape of the MZR at low mass not entirely representative of main sequence galaxies. Clearly, a thorough study using the deepest spectroscopy available for a statistical significant complete sample of SFDGs is much needed to test this hypothesis. Forthcoming analysis of VUDS galaxies at $z<1$ using the complete database will enable us to scrutinize in detail the underexplored low-mass universe at $z<1$.
14
3
1403.3692
We report the discovery of 31 low-luminosity (-14.5 ≳ M<SUB>AB</SUB>(B) ≳ -18.8), extreme emission line galaxies (EELGs) at 0.2 ≲ z ≲ 0.9 identified by their unusually high rest-frame equivalent widths (100 ≤ EW[O iii] ≤ 1700 Å) as part of the VIMOS Ultra Deep Survey (VUDS). VIMOS optical spectra of unprecedented sensitivity (I<SUB>AB</SUB> ~ 25 mag) along with multiwavelength photometry and HST imaging are used to investigate spectrophotometric properties of this unique sample and to explore, for the first time, the very low stellar mass end (M<SUB>⋆</SUB> ≲ 10<SUP>8</SUP>M<SUB>⊙</SUB>) of the luminosity-metallicity (LZR) and mass-metallicity (MZR) relations at z &lt; 1. Characterized by their extreme compactness (R<SUB>50</SUB> &lt; 1 kpc), low stellar mass and enhanced specific star formation rates (sSFR = SFR/M<SUB>⋆</SUB> ~ 10<SUP>-9</SUP>-10<SUP>-7</SUP> yr<SUP>-1</SUP>), the VUDS EELGs are blue dwarf galaxies likely experiencing the first stages of a vigorous galaxy-wide starburst. Using T<SUB>e</SUB>-sensitive direct and strong-line methods, we find that VUDS EELGs are low-metallicity (7.5 ≲ 12 + log (O/H) ≲ 8.3) galaxies with high ionization conditions (log (q<SUB>ion</SUB>) ≳ 8 cm s<SUP>-1</SUP>), including at least three EELGs showing Heiiλ 4686 Å emission and four extremely metal-poor (≲10% solar) galaxies. The LZR and MZR followed by VUDS EELGs show relatively large scatter, being broadly consistent with the extrapolation toward low luminosity and mass from previous studies at similar redshift. However, we find evidence that galaxies with younger and more vigorous star formation - as characterized by their larger EWs, ionization and sSFR - tend to be more metal poor at a given stellar mass. <P />Based on data obtained with the European Southern Observatory Very Large Telescope, Paranal, Chile, under Large Program 185.A-0791.Figure A.1 is available in electronic form at <A href="http://www.aanda.org/10.1051/0004-6361/201423816/olm">http://www.aanda.org</A>Tables 1 and 2 are only available at the CDS via anonymous ftp to <A href="http://cdsarc.u-strasbg.fr">http://cdsarc.u-strasbg.fr</A> (ftp://130.79.128.5) or via <A href="http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/568/L8">http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/568/L8</A>
false
[ "VUDS EELGs", "blue dwarf galaxies", "extreme emission line galaxies", "high ionization conditions", "low luminosity", "similar redshift", "VUDS", "anonymous ftp", "mass", "more metal", "first", "HST imaging", "low stellar mass and enhanced specific star formation rates", "previous studies", "multiwavelength photometry", "z", "electronic form", "spectrophotometric properties", "a given stellar mass", "ionization" ]
12.659989
7.016151
-1
484492
[ "Line, Michael R.", "Fortney, Jonathan J.", "Marley, Mark S.", "Sorahana, Satoko" ]
2014ApJ...793...33L
[ "A Data-driven Approach for Retrieving Temperatures and Abundances in Brown Dwarf Atmospheres" ]
39
[ "Department of Astronomy and Astrophysics, University of California, Santa Cruz, CA 95064, USA", "Department of Astronomy and Astrophysics, University of California, Santa Cruz, CA 95064, USA", "NASA Ames Research Center, Mail Stop 245-3; Moffett Field, CA 94035, USA", "Department of Physics, Nagoya University, Nagoya, Aichi 464-8602, Japan" ]
[ "2015ARA&A..53..279M", "2015ApJ...799...37L", "2015ApJ...807..183L", "2015MNRAS.453.3414A", "2015PASP..127..941C", "2016ApJ...817...17G", "2016ApJ...829...66M", "2016ApJ...833..120R", "2016PASP..128j2001B", "2016arXiv161203174M", "2017ApJ...848...83L", "2017AstRv..13....1B", "2017MNRAS.464.4247W", "2017MNRAS.470.1177B", "2017PASP..129c4401N", "2017arXiv170802821A", "2018A&A...611A...8S", "2018arXiv180408149F", "2018haex.bookE.101B", "2019ApJS..240...19K", "2020A&A...640A.131M", "2020AJ....159....6O", "2020AJ....159..289L", "2020MNRAS.497.5136P", "2021A&A...653A..73S", "2021ApJ...913..114W", "2021ApJ...921L...4L", "2021ApJ...922...26T", "2021ApJ...923...19G", "2022ApJ...936...44Z", "2023A&A...673A.125S", "2023AJ....166...86S", "2023AJ....166..192A", "2023ApJ...944..138V", "2023ApJ...953..170H", "2023MNRAS.522.5062B", "2023MNRAS.525.1375W", "2024AJ....167..237L", "2024arXiv240606493K" ]
[ "astronomy" ]
7
[ "brown dwarfs", "methods: statistical", "radiative transfer", "stars: atmospheres", "stars: individual: Gl 570D", "Astrophysics - Solar and Stellar Astrophysics", "Astrophysics - Earth and Planetary Astrophysics" ]
[ "1976RvGSP..14..609R", "1993ApJ...406..158B", "1996ApJ...465L.123A", "1996Sci...272.1919M", "1997A&A...323L..49P", "1997ARA&A..35..137A", "1997Sci...276..108Y", "1998Icar..135..501C", "1999ApJ...519L..85G", "2000ApJ...531L..57B", "2000ApJ...541..374S", "2000imas.book.....R", "2001ApJ...556..357A", "2001ApJ...556..373G", "2002ApJ...564..452L", "2003ApJ...594..510B", "2003PASP..115..362R", "2004AdSpR..34..700C", "2004SPIE.5487...62H", "2005AJ....129.1706F", "2006ApJ...637.1067B", "2006ApJ...640.1063B", "2006ApJ...647..552S", "2006ApJ...648..614C", "2006ApJ...651..502P", "2007ApJ...657..494B", "2007ApJ...669.1248H", "2007Icar..188...47N", "2007Icar..189..457F", "2008ApJ...678.1372C", "2008ApJ...689.1327S", "2008ApJS..174..504F", "2008JQSRT.109.1136I", "2009ApJ...695..844G", "2009ApJ...702..154S", "2009ApJ...707...24M", "2009JQSRT.110..533R", "2010ApJ...722..682Y", "2010ApJS..186...63R", "2011ApJ...734...73T", "2011ApJ...737...15M", "2011ApJ...738...72V", "2011ApJ...743..191M", "2011ApJ...743L..16O", "2011Icar..214..606G", "2011MNRAS.413.1828Y", "2012A&A...540A..85P", "2012ApJ...747L..27F", "2012ApJ...749...93L", "2012ApJ...750...74S", "2012ApJ...753..100B", "2012ApJ...760..151S", "2012MNRAS.420..170L", "2013A&A...552A..73N", "2013ApJ...775..137L", "2013ApJ...778...97L", "2013ApJ...778..132T", "2013MNRAS.434.2616B", "2014ApJ...783...70L", "2014Life....4..142H", "2014MNRAS.440.3675S" ]
[ "10.1088/0004-637X/793/1/33", "10.48550/arXiv.1403.6412" ]
1403
1403.6412_arXiv.txt
Unlike most stars, with photospheres that predominantly emit over a limited range of altitudes, the molecular opacity-dominated atmospheres of brown dwarfs have large variations in opacity with wavelength, allowing flux from very different atmospheric depths to emerge over various spectral ranges. This property--which brown dwarfs share with planetary atmospheres--allows information to be extracted from a range of altitudes and conditions, if spectral measurements are available from a sufficiently broad swath of wavelengths. Historically, by comparing models to observations, observed spectra have been interpreted to discern brown dwarf masses, formation modes, evolution, and the processes at work in their atmospheres (Burrows et al. 1993; Allard et al. 1996; Allard et al. 2001; Marley et al. 1996; Saumon et al. 2000; Geballe et al. 2001; Burrows et al. 2006; Hubeny \& Burrows 2007; Burgasser et al. 2007; Cushing et al. 2008; Stephens et al. 2009; Rice et al. 2010; Yamamura et al. 2010; Patience et al. 2012). These processes include vertical mixing, disequilibrium chemistry, global dynamics, and cloud formation and sedimentation. Currently, spectra of brown dwarfs are generally interpreted through comparisons of observed spectra with pre-computed grids of model atmospheres. Such models self-consistently solve for the temperature structure in radiative-convective equilibrium and molecular abundances in the atmosphere with few free parameters, typically $\log g$ and effective temperature, cloud parameterizations, and in some cases an eddy mixing coefficient to accommodate non-equilibrium chemistry due to vertical mixing. The combination of $\log g$ and effective temperature that provide a best fit, and the subsequent atmospheric structure, are taken to be the solution for a particular brown dwarf's atmosphere. Within the framework parameter grid search or Monte-Carlo methods are sometimes implemented in order to estimate the uncertainties in these parameters (e.g., Cushing et al. 2008; Rice et al. 2010). While this grid-based comparison approach has offered considerable insight into interpretation of the spectra, these methods are constrained by various assumptions which do not easily allow for a full exploration of brown dwarf parameter space. For instance, most self-consistent grid models assume thermochemical equilibrium, which we now know is generally not the case, especially with molecular species like NH$_3$/N$_2$ and CH$_4$/CO which can be driven strongly out of equilibrium due to vertical mixing (e.g., Saumon et al. 2006; Griffith \& Yelle 1999) . Within the self consistent grid model frame work this is remedied by some via the inclusion 1D vertical mixing prescription parameterizied with an eddy diffusivity parameter, but such an approach is often reliant upon the choosing the correct rate limiting steps (e.g., see Visscher \& Moses 2011; Moses et al. 2011). Furthermore, most of these models generally assume solar elemental abundances, such that modeling investigations of metal-poor and metal-rich atmospheres, or deviations from solar-like C/N/O ratios have generally been lacking (although see Allard 1997; Tsuji et al. 2011; Burrows et al. 2006; Saumon \& Marley 2008). The underlying causes for deviations of best-fitting models from data have rarely been explored, even though such deviations record shortcomings in the underlying model assumptions or underlying physical data. The vast majority of these models also assume 1D radiative-convective equilibrium in a static atmosphere, leaving one unable to explore departures from these temperature structures arising from dynamical transport of heat, disequilibrium chemistry, latent heat release due to condensation, or other phenomena. Furthermore, given the widespread observational evidence for variability in the emitted spectra of brown dwarfs, new modeling approaches that relax assumptions inherent in previous models are now needed. The goal of this investigation is to introduce a new approach for the interpretation of brown dwarf spectra. Powerful techniques to directly invert measured spectra into constraints on molecular abundances and atmospheric temperature structure have been widely used within the Earth sciences (Rodgers 1976, Towmey 1996, Rodgers 2000, Crisp et al. 2004) and solar system planets (Conrath et al. 1998, Irwin et al. 2008, Nixon et al. 2007, Fletcher et al. 2007, Greathouse et al. 2011) and recently exoplanet atmosphere inference (Lee et al. 2012; 2013 Line et al. 2012; Barstow et al. 2013; Line et al. 2013a). These atmospheric retrieval approaches are primarily data driven and free from many of the aforementioned assumptions, and naturally allow for a wide exploration of brown dwarf atmosphere parameter space. The atmospheric temperature structure and molecular abundances are ``retrieved" through the iterative calculation of many tens to thousands of model spectra. Each spectrum is generated with a unique temperature structure and variations on the molecular abundances. An assessment is made regarding the goodness of fit of each model. Since an extremely large phase space of temperature structures and abundances are probed, the end product is a range of models (or an analytic estimate of that range) that fit the observed spectra, along with a statistical assessment of the goodness of fit from a wide range of models. In the current paper we apply the above well established retrieval methodologies to brown dwarf spectra in order to illustrate the veracity of the approach. We first retrieve the thermal structure and atmospheric abundances of a model and then turn to the well-studied T dwarf Gl 570D, which has previous been the target of extensive observational (Burgasser 2000;2003;2006; Leggett et al. 2002; Cushing et al. 2006; Patten et al. 2006; Geballe et al. 2009) and modeling campaigns (Geballe et al. 2000; Saumon et al. 2006). The paper is organized as follows: In \S\ref{sec:Methods} we introduce the retrieval methodology. A synthetic example is shown in \S\ref{sec:Synthetic} and our initial results for Gl570D are in \S\ref{sec:Gl570}. Finally, in \S\ref{sec:Conclusions} future research directions are discussed.
\label{sec:Conclusions} We have for the first time applied well proven, data driven, temperature and abundance retrieval techniques to a brown dwarf spectrum free from many of the physical assumptions present in grid models. These approaches allow an unbiased determination of the temperatures and abundances in brown dwarf atmospheres. We first demonstrated that the optimal estimation retrieval approach is a powerful atmospheric inference tool in the presence of gaussian uncertainties. With a model that is will matched to the data and minimal systematic uncertainties for typical observational conditions, abundances in brown dwarf atmospheres can be determined to within a few tens of percent, compared with orders-of-magnitude for exoplanet data. Furthermore, the full temperature profile can be constrained at most pressure levels to better than 50 K. We then applied our approach to the well studied brown dwarf, Gl 570D. For the first time, by combining SpeX, AKARI, and Spitzer IRS data we were able to constrain the abundances of water, methane, carbon monoxide, carbon dioxide, and ammonia as well as temperature structure, gravity and radius (and hence mass). Real data is plagued with systematic uncertainties due to photometric calibration errors and missing forward model physics. We were able to at least accommodate for the photometric uncertainties with the bootstrap Monte Carlo approach. In lieu of these systematic uncertainties we were able to constrain the molecular mixing ratios to better than order-of-magnitude precisions, the temperature profile to within $\sim$100 K at most pressure levels, the effective temperature to within $\sim$50 K, and the mass to within a factor of two. We found that our results are fairly consistent with those of Saumon et al. (2006) with the largest difference being a lower retrieved ammonia abundance. As with any data-model comparison method, our approach does have some shortcomings and limitations. First and foremost, we do not test to see if the derived thermal profile with the derived abundances are in radiative equilibrium in the upper part of the atmosphere. It may be that our solution would be expected to relax rapidly to a different thermal state. We simply allow the data to guide the solutions rather than our preconceived notions on how the atmosphere {\em should} behave. Therefore within this weakness lies our potential strength in the sense that we are not bound by these assumptions; we can test whether or not such assumptions are correct. For instance, by not assuming radiative equilibrium we allow for other possible heating mechanisms, such as gravity wave breaking (Young et al. 1997) or non-LTE heating processes (Sorahana et al. 2014), not commonly accounted for within the standard grid modeling framework. Another weakness, with the current forward model at least, is the assumption of abundances that are uniform with height. This assumption and whether or not the data justifies including more complicated abundance profiles can be tested in the future. Additionally, we have not included clouds, since for this object we are comparing results from previous cloud-free grid models (e.g., Saumon et al. 2006; Geballe et al. 2009). We may also be limited by the accuracy of the cross section databases. These short comings fall under the more general category of``missing model physics" which when properly accounted for (a task for future investigations) may result in an overall increase the uncertainties on the desired model parameters. Another potential issue is the applicability of our approach to lower quality data. In these scenarios the gaussian posterior assumption used in the optimal estimation approach may not be valid. Line et al. (2013) explore in full detail when this assumption is valid, and when more sophisticated techniques (e.g., bootstrap monte carlo and Markov chain Monte Carlo) are necessary. Finally, there is a continuum of approaches with which one may use in the atmospheric retrieval problem. On one end, a simple forward model with few limiting assumptions can be used, as was done in this investigation. On the other end is the self-consistent grid modeling approach where by many assumptions are made to reduce the problem to just a few simple parameters. It is worth while in future investigations to explore this continuum and to understand how differing physical assumptions can change the results. Our philosophy is to start with the minimal number of assumptions and build in layers of sophistication as we go. There are many questions that we can begin to ask and potentially answer with our retrieval approach. What is the atmospheric temperature structure of Brown Dwarfs? Are they in radiative equilibrium as current theory suggests? Can deviations from radiative equilibrium be a potential diagnostic for dynamical processes? What are the compositions of the brown dwarfs? Do brown dwarfs have non-solar elemental ratios? In other words, is the metallicity enhancement different for each element? Does composition vary with altitude? How do they deviate from thermochemical equilibrium? What is the vigor of vertical mixing? Are the observations even sensitive to vertical variations in the abundance profiles? How will clouds impact the retrievals? Can we determine the cloud opacities and cloud levels? Is there missing physics in the self consistent models and can we aid in identification in that missing physics? Are the observations sensitive to thermal variations in the temperature profile as a function of time? Are the elemental compositional differences between different objects? Does this inform us on their formation environments? The unprecedented quality of brown dwarf data will allow us to begin to address many of these tantalizing questions.
14
3
1403.6412
Brown dwarf spectra contain a wealth of information about their molecular abundances, temperature structure, and gravity. We present a new data driven retrieval approach, previously used in planetary atmosphere studies, to extract the molecular abundances and temperature structure from brown dwarf spectra. The approach makes few a priori physical assumptions about the state of the atmosphere. The feasibility of the approach is first demonstrated on a synthetic brown dwarf spectrum. Given typical spectral resolutions, wavelength coverage, and noise, property precisions of tens of percent can be obtained for the molecular abundances and tens to hundreds of K on the temperature profile. The technique is then applied to the well-studied brown dwarf, Gl 570D. From this spectral retrieval, the spectroscopic radius is constrained to be 0.75-0.83 R <SUB>J</SUB>, log (g) to be 5.13-5.46, and T <SUB>eff</SUB> to be between 804 and 849 K. Estimates for the range of abundances and allowed temperature profiles are also derived. The results from our retrieval approach are in agreement with the self-consistent grid modeling results of Saumon et al. This new approach will allow us to address issues of compositional differences between brown dwarfs and possibly their formation environments, disequilibrium chemistry, and missing physics in current grid modeling approaches as well as a many other issues.
false
[ "Brown dwarf spectra", "brown dwarf spectra", "brown dwarfs", "temperature profiles", "temperature structure", "current grid modeling approaches", "abundances", "planetary atmosphere studies", "Saumon et al", "compositional differences", "missing physics", "a synthetic brown dwarf spectrum", "disequilibrium chemistry", "the molecular abundances", "their molecular abundances", "issues", "tens", "the temperature profile", "Gl 570D.", "property precisions" ]
7.193036
12.648544
-1
812633
[ "Avelino, P. P.", "Bolejko, K.", "Lewis, G. F." ]
2014PhRvD..89j3004A
[ "Nonlinear Chaplygin gas cosmologies" ]
31
[ "Centro de Astrofísica da Universidade do Porto, Rua das Estrelas, 4150-762 Porto, Portugal; Departamento de Física e Astronomia, Faculdade de Ciências da Universidade do Porto, Rua do Campo Alegre 687, 4169-007 Porto, Portugal; Sydney Institute for Astronomy, School of Physics, A28, The University of Sydney, Camperdown, New South Wales 2006, Australia", "Sydney Institute for Astronomy, School of Physics, A28, The University of Sydney, Camperdown, New South Wales 2006, Australia", "Sydney Institute for Astronomy, School of Physics, A28, The University of Sydney, Camperdown, New South Wales 2006, Australia" ]
[ "2015JCAP...05..020C", "2015JCAP...10..056C", "2015MNRAS.453.2250E", "2015PhRvD..91h3508A", "2015PhRvD..92j3511K", "2016EPJC...76..567S", "2016IJMPD..2550031J", "2016JCAP...01..040B", "2016MPLA...3150202C", "2016arXiv160901477S", "2017CaJPh..95.1068R", "2017IJTP...56.3297S", "2017PhLB..770..213F", "2017arXiv170100748V", "2018EPJC...78..495S", "2018PhRvD..98d3515F", "2019A&A...623A..28L", "2019CaJPh..97..286B", "2019JCAP...11..044Y", "2019PhRvD..99b3532C", "2020IJGMM..1750194H", "2020PhRvD.102f3525F", "2021IJGMM..1850095H", "2022CQGra..39s5007Y", "2022EPJC...82..455E", "2022NewA...9501823V", "2022PhRvD.106f3525C", "2022PhRvD.106h3524A", "2022PhyS...97d5006S", "2023Symm...15.1332S", "2024arXiv240601705C" ]
[ "astronomy", "physics" ]
6
[ "95.36.+x", "95.35.+d", "98.80.Es", "98.80.Cq", "Dark energy", "Dark matter", "Observational cosmology", "Particle-theory and field-theory models of the early Universe", "Astrophysics - Cosmology and Nongalactic Astrophysics", "General Relativity and Quantum Cosmology" ]
[ "2000PhRvD..62h5023O", "2001PhLB..511..265K", "2002PhLB..535...17B", "2002PhLB..537....1G", "2002PhRvD..66b1301P", "2002PhRvD..66d3507B", "2003JCAP...07..005A", "2003JCAP...09..002A", "2003PhLB..555....1M", "2003PhLB..575..172B", "2003PhRvD..67b3515D", "2003PhRvD..67d3514A", "2003PhRvD..67f3504B", "2003PhRvD..67j1301B", "2003PhRvD..68d3504B", "2003PhRvD..68f1302R", "2004A&A...417..847D", "2004ApJ...604....1F", "2004JCAP...11..008B", "2004MNRAS.353..329B", "2004PhRvD..69d1301A", "2004PhRvD..69h3501C", "2004PhRvD..69l3517C", "2004PhRvD..69l3524S", "2005ApJ...622...28B", "2005CQGra..22.4311Z", "2005JCAP...06..007G", "2006IJMPD..15.1753C", "2006PhRvD..74d3528C", "2007JCAP...03..015W", "2007JPhA...40.6877B", "2007MNRAS.376.1169B", "2007PhLB..644...16W", "2008JCAP...02..016G", "2008PhLB..662...87L", "2008PhRvD..77f3515A", "2008PhRvD..77j1302A", "2008PhRvD..78b3505B", "2010AdAst2010E..78B", "2010CQGra..27q5013R", "2010JCAP...03..025X", "2010PhRvD..81h7303F", "2010PhRvD..82f3534A", "2011CQGra..28p4002B", "2011CoTPh..56..525L", "2011EPJC...71.1773F", "2011GReGr..43..819L", "2011PhLB..699...10A", "2012ARNPS..62...57B", "2012ApJ...746...85S", "2012EPJC...72.1931X", "2012MNRAS.427.3435A", "2012PhLB..714....6A", "2012PhR...513....1C", "2012PhRvD..85d3007C", "2012PhRvD..86j3518P", "2013PhRvD..87d3527D", "2013PhRvD..87h3503W" ]
[ "10.1103/PhysRevD.89.103004", "10.48550/arXiv.1403.1718" ]
1403
1403.1718_arXiv.txt
Cosmological observations provide overwhelming evidence for Dark Energy (DE) \cite{Suzuki:2011hu,Anderson:2012sa,Parkinson:2012vd,Ade:2013zuv}. This of course is a subject to several assumptions such as that Einstein's General Relativity provides an accurate description of gravity on cosmological scales, and in addition that the Friedmann--Lema\^itre--Robertson--Walker models adequately describe our Universe \cite{Clifton:2011jh,Bolejko:2011jc,Buchert:2011sx}. Although the simplest DE candidate, a cosmological constant $\Lambda$, is perfectly consistent with current cosmological observations, there is presently no satisfactory explanation for the tiny DE density (over $120$ orders of magnitude smaller than the Planck density) which, nevertheless, appears to account for about $70 \%$ of the total energy density of the Universe at the present time. Hence, despite the simplicity of the cosmological constant, dynamical DE models are arguably better motivated from a theoretical point of view (see, for example, \cite{Copeland:2006wr,Li:2011sd} and references therein). While DE might explain the observed dynamics of the Universe on cosmological scales, a nonrelativistic dark matter (DM) component is required in the standard cosmological model to account for the observed clustering of large scale structures. The simplest model which incorporates the DM and DE components is the so-called $\Lambda$CDM model, in which the DE and DM roles are played by a cosmological constant $\Lambda$ and Cold Dark Matter (CDM) particles with a negligible free streaming length. The energy components of this model can either be taken as a DE fluid with $p_{\rm DE}=-\rho_{\rm DE}=-\rho_\Lambda$ and a DM fluid with $p_{\rm DM}=0$ or a single Unified DE (UDE) component with $p_{DE}=-\rho_\Lambda$ and arbitrary $\rho > \rho_\Lambda$ (here $\rho$ and $p$ represent the density and pressure, respectively). Hence, the $\Lambda$CDM scenario can be regarded as the simplest example of a UDE realization where the role of DM and DE are played by the same dark fluid \cite{Avelino:2003cf}. Various other interesting candidates for the unification of DM and DE have been proposed in the literature, including the Chaplygin gas and its variations \cite{Kamenshchik:2001cp,Bilic:2001cg,Bento:2002ps}, tachyon field models \cite{Gibbons:2002md,Padmanabhan:2002cp,Bagla:2002yn,Chimento:2003ta,Calcagni:2006ge,Avelino:2010qn,Avelino:2011ey}, and a large variety of Interacting Dark Energy (IDE) models \cite{Bassett:2002fe,Farrar:2003uw,Gumjudpai:2005ry,Boehmer:2008av,Clemson:2011an,Avelino:2012tc}. In this paper we shall focus on UDE models in which the UDE fluid is described by a perfect fluid with an isentropic Equation of State (EoS) $p=p(\rho) = w(\rho) \rho$, where $w(\rho)$ is the EoS parameter of the DE, but most of our results will also apply to other IDE models in the strong coupling regime. Despite the very different parameterizations available for $w(\rho)$, all UDE models are characterized by an EoS, satisfying two important properties: i) if the density is much larger than the average density of the Universe at the present time, then $w \sim 0$ (and $c_s \sim 0$, where $c_s$ is the sound speed); ii) if the density is close to the current average density of the Universe then the EoS parameter of the UDE fluid is close to $-1$. A representative example of a family of isentropic UDE models is the Generalized Chaplygin Gas (GCG), characterized by the EoS $p=-A/\rho^{\alpha}$ where $A > 0$ and $0 \le \alpha \le 1$ are constants. Isentropic UDE models have been claimed to be essentially ruled out due to the late time oscillations of the matter power spectrum inconsistent with observations, except for a small region of parameter space close to the standard $\Lambda$CDM model \cite{Sandvik:2002jz} (for $\alpha < 0$ linear theory would instead predict an exponential blowup of the matter power spectrum). Although the inclusion of baryons in the analysis does lead to less stringent bounds on the GCG parameter $\alpha$ \cite{Beca:2003an}, linear isentropic UDE models have been shown to be tightly constrained by cosmological observations \cite{Alcaniz:2002yt,Dev:2002qa,Makler:2002jv,Bento:2003we,Amendola:2003bz,Cunha:2003vg,Dev:2003cx,Bertolami:2004ic,Biesiada:2004td,Wu:2006pe,Wu:2007bv,Gorini:2007ta,Lu:2008zzb,Fabris:2010yh,Xu:2010zzb,Lu:2010zzj,Fabris:2011wk,Xu:2012ca,Wang:2013qy} (see also \cite{Reis:2003mw,Zimdahl:2005ir,Bilic:2006cp,Bertacca:2010ct} for a discussion of nonisentropic UDE models). The effect of small scale nonlinearities has been recognized as having a strong potential impact on the large scale evolution of the Universe, in particular in the context of UDE scenarios, \cite{Avelino:2003ig,Bilic:2003cv,Beca:2005gc,Avelino:2007tu,Avelino:2008cu,Roy:2009cj,DelPopolo:2013bpa}. However, it has been argued that, in the case of the Chaplygin gas, nonlinear effects would be too small to significantly affect the above linear results \cite{Bilic:2003cv} (see also \cite{Avelino:2003ig}). This conclusion relied on the assumption of a constant spectral index of scalar gaussian fluctuations as well as an EoS for the Chaplygin gas whose form remains unchanged both at large densities and small scales. These are very strong assumptions, given the relatively weak constraints on the scalar spectral index on nonlinear scales (wavenumbers $k \gsim 3 \, {\rm Mpc}^{-1}$) and the expectation that the isentropic perfect fluid approximation might break at sufficiently large densities or small scales. In this paper we relax these assumptions and model the effect of the small scale nonlinearities on the background evolution of the Universe using a single parameter $\epsilon$, representing the fraction of the UDE density which has become nonlinear due to the gravitational collapse into UDE halos. We show that, for $\epsilon$ close to unity, the linear theory results no longer hold and that the backreaction of the small scale structures on the large scale evolution of the Universe may render the Chaplygin gas model virtually indistinguishable from the $\Lambda$CDM scenario for all possible values of the GCG parameter $\alpha$. In this paper we shall use units where $8 \pi G/3=1$.
} In this paper we have parametrised the effect of UDE nonlinear clustering on the dynamics of the Universe. We have shown that cosmological scenarios in which the DM and DE roles are played by a single UDE fluid may be reconciled with the latest observational results, provided there is a high level of nonlinear clustering of the UDE component. Although we have focused on the GCG as a concrete example, our main results are expected to hold in general for UDE models.
14
3
1403.1718
We study the nonlinear regime of unified dark energy models, using generalized Chaplygin gas cosmologies as a representative example, and introduce a new parameter characterizing the level of small scale clustering in these scenarios. We show that viable generalized Chaplygin gas cosmologies, consistent with the most recent observational constraints, may be constructed for any value of the generalized Chaplygin gas parameter by considering models with a sufficiently high level of nonlinear clustering.
false
[ "generalized Chaplygin gas cosmologies", "viable generalized Chaplygin gas cosmologies", "unified dark energy models", "nonlinear clustering", "the generalized Chaplygin gas parameter", "Chaplygin", "models", "a new parameter", "these scenarios", "a representative example", "the nonlinear regime", "the most recent observational constraints", "the level", "a sufficiently high level", "any value", "We" ]
11.483136
0.405939
-1
788371
[ "Kato, Taichi", "Osaki, Yoji" ]
2014PASJ...66L...5K
[ "GALEX J194419.33+491257.0: An unusually active SU UMa-type dwarf nova with a very short orbital period in the Kepler data" ]
14
[ "Department of Astronomy, Kyoto University, Kitashirakawa-Oiwake-cho, Sakyo-ku, Kyoto 606-8502", "Department of Astronomy, School of Science, The University of Tokyo, 7-3-1 Hongo, Bunkyo-ku, Tokyo 113-0033" ]
[ "2014ApJ...790...28R", "2014PASJ...66...67O", "2014PASJ...66...90K", "2015AJ....149...67B", "2015ARA&A..53..247M", "2015PASJ...67..105K", "2016AJ....151..159S", "2016PASJ...68..107K", "2021MNRAS.505.2051E", "2022ApJS..260...20F", "2022Natur.610..467B", "2022arXiv220102945K", "2023ApJ...950..139G", "2023ApJ...953L...1B" ]
[ "astronomy" ]
6
[ "accretion", "accretion disks", "novae", "cataclysmic variables", "stars: dwarf novae", "stars: individual (GALEX J194419.33+491257.0)", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1989PASJ...41.1005O", "1995PASJ...47..163K", "1995PASP..107..443R", "1996A&A...311..889A", "1999dicb.conf...45K", "2001PASJ...53.1191K", "2002ApJ...567L..49T", "2002PASJ...54L..15U", "2003MNRAS.340.1214P", "2004MNRAS.349..181N", "2005ApJ...619L...1M", "2006PASJ...58..143I", "2009PASJ...61S.395K", "2010ApJ...713L..79K", "2010Sci...327..977B", "2011A&A...536A..42Z", "2011ApJS..194...28K", "2011MNRAS.415.2025S", "2012MNRAS.419.2900F", "2012MNRAS.422.1219B", "2013MNRAS.434.1902P", "2013PASJ...65...95O", "2013PASJ...65...97K", "2013PASJ...65..115K", "2013PASP..125..506T", "2014PASJ...66...15O" ]
[ "10.1093/pasj/psu025", "10.48550/arXiv.1403.0308" ]
1403
1403.0308_arXiv.txt
The Kepler mission (\cite{bor10Keplerfirst}; \cite{Kepler}), which was aimed to detect extrasolar planets, has provided unprecedentedly sampled data on several cataclysmic variables (CVs). This satellite also recorded previously unknown CVs as by-products of the main target stars. The best documented example has been the background dwarf nova of KIC 4378554 (\cite{bar12j1939}; \cite{kat13j1939v585lyrv516lyr}). In addition to this object, the group Planet Hunters \citep{fis12PlanetHunters} detected several candidate background CVs.\footnote{ $<$http://talk.planethunters.org/objects/APH51255246/\\ discussions/DPH101e5xe$>$. } We studied one of these background dwarf novae, the one in the field of KIC 11412044 (hereafter J1944). This object was discovered by the Planet Hunters group as a background SU UMa-type dwarf nova of KIC 11412044, in which superoutbursts and frequent normal outbursts were recognized. \footnote{ $<$http://keplerlightcurves.blogspot.jp/2012/07/\\ dwarf-novae-candidates-at-planet.html$>$. } Since it was bright enough and it was frequently included in the aperture mask of KIC 11412044, the outburst behavior can be immediately recognized in Kepler {\tt SAP\_FLUX} light curve of KIC 11412044.
14
3
1403.0308
We studied a background dwarf nova of KIC 11412044 in the Kepler public data and identified it with GALEX J194419.33+491257.0. This object turned out to be a very active SU UMa-type dwarf nova that has a mean supercycle of ∼ 150 d and frequent normal outbursts with intervals of 4-10 d. The object showed a strong persistent signal of the orbital variation with a period of 0.0528164(4) d (76.06 min) and superhumps with a typical period of 0.0548 d during its superoutbursts. Most of the superoutbursts were accompanied by a precursor outburst. All these features are unusual for this very short orbital period. We succeeded in detecting an evolving stage of superhumps (stage A superhumps) and obtained a mass ratio of 0.141(2), which is unusually high for this orbital period. We suggest that the unusual outburst properties are a result of this high mass ratio. We suspect that this object is a member of the recently recognized class of cataclysmic variables (CVs) with a stripped core evolved secondary which are evolving toward AM CVn-type CVs. The present determination of the mass ratio by using stage A superhumps is the first case in such systems.
false
[ "frequent normal outbursts", "d.", "such systems", "intervals", "superhumps", "stage", "CVs", "AM CVn-type CVs", "0.0528164(4", "a precursor outburst", "this orbital period", "a typical period", "nova", "cataclysmic variables", "an evolving stage", "this high mass ratio", "GALEX", "stage A superhumps", "first", "this very short orbital period" ]
5.128291
9.86785
26
812604
[ "Miranda, Vinícius", "Hu, Wayne", "Adshead, Peter" ]
2014PhRvD..89j1302M
[ "Steps to reconcile inflationary tensor and scalar spectra" ]
54
[ "Department of Astronomy and Astrophysics, University of Chicago, Chicago, Illinois 60637, USA; The Capes Foundation, Ministry of Education of Brazil, Brasília DF 70359-970, Brazil", "Department of Astronomy and Astrophysics, University of Chicago, Chicago, Illinois 60637, USA; Kavli Institute for Cosmological Physics, Enrico Fermi Institute, University of Chicago, Chicago, Illinois 60637, USA", "Department of Physics, University of Illinois at Urbana-Champaign, Urbana, Illinois 61801, USA" ]
[ "2014EPJC...74.2954Z", "2014JCAP...06..061H", "2014JCAP...07..020C", "2014JCAP...07..033C", "2014JCAP...08..036M", "2014JCAP...08..043K", "2014JCAP...08..044E", "2014JCAP...08..048H", "2014JCAP...08..053A", "2014JCAP...10..025H", "2014JCAP...10..044Z", "2014JCAP...11..012M", "2014JCAP...12..019B", "2014JHEP...07..044H", "2014JHEP...07..074H", "2014MNRAS.442..670C", "2014PhLB..734...41G", "2014PhLB..734..167C", "2014PhLB..734..345C", "2014PhLB..737..191H", "2014PhRvD..89l3503H", "2014PhRvD..90b3511A", "2014PhRvD..90b3537W", "2014PhRvD..90b3544H", "2014PhRvD..90f3007B", "2014PhRvD..90f3525F", "2014PhRvD..90h3503D", "2014PhRvD..90h3515D", "2014PhRvD..90j3531L", "2014PhRvD..90l3524H", "2014PhRvL.112s1301D", "2014PhRvL.112s1303B", "2014PhRvL.112y1301X", "2014PhRvL.113c1301S", "2014PhRvL.113g1301H", "2014SCPMA..57.1414C", "2014SCPMA..57.1455L", "2014arXiv1404.0238L", "2014arXiv1404.3804X", "2015IJMPD..2430023C", "2015JCAP...01..010E", "2015JCAP...01..040K", "2015JCAP...08..035B", "2015PhRvD..91f3509G", "2015PhRvD..91f3514M", "2015PhRvD..91f3521K", "2015PhRvD..92f3511C", "2016JCAP...06..027C", "2016PhRvD..94b3526B", "2017PDU....17...38G", "2017PhRvD..95l3534G", "2019arXiv190609299R", "2022JPhCS2372a2002N", "2024IJMPD..3350002C" ]
[ "astronomy", "physics" ]
3
[ "98.80.Cq", "98.70.Vc", "Particle-theory and field-theory models of the early Universe", "Background radiations", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
[ "1975CoPhC..10..343J", "2000ApJ...538..473L", "2001PhRvD..64l3514A", "2002PhRvD..65j3508S", "2003ApJS..148..213P", "2004PhRvD..70j3505S", "2004PhRvD..70l3505A", "2006JCAP...09..010E", "2007JCAP...06..023C", "2008JCAP...04..010C", "2009PhRvD..79j3519M", "2010PhRvD..81b3518D", "2011PhRvD..84b7303H", "2011PhRvD..84d3519A", "2012JCAP...04..027H", "2012PhRvD..85b3531A", "2012PhRvD..85h3520P", "2012PhRvD..85j3531A", "2012PhRvD..86f3529M", "2012PhRvD..86l1301A", "2013PhRvD..88b3507A", "2014A&A...571A..22P", "2014JCAP...07..014C", "2014PhRvD..89h3529M", "2014PhRvL.112x1101B" ]
[ "10.1103/PhysRevD.89.101302", "10.48550/arXiv.1403.5231" ]
1403
1403.5231_arXiv.txt
The recent BICEP2 measurement of a tensor-scalar ratio $r=0.2^{+0.07}_{-0.05}$ from degree scale B-mode polarization of the cosmic-microwave background (CMB) \cite{Ade:2014xna} is in ``moderately-strong" tension with slow-roll inflation models that predict scale-free, albeit slightly tilted ($1-n_s \ll 1$) power-law power spectra. This tension is due to the implied excess in the temperature spectrum at low multipoles which is not observed and restricts $r_{0.002}< 0.11$ (95\% CL) in this context \cite{Ade:2013uln}. These findings can be reconciled in the single-field inflationary paradigm by introducing a scale into the scalar power spectra to suppress power on these large-angular scales. For example a large running of tilt, $dn_s/d\ln k \sim -0.02$, is possible as a compromise \cite{Ade:2014xna}. Here the scale introduced is associated with the scalar spectrum transiently passing through a scale-invariant slope near observed scales. However, such a large running is uncomfortable in the simplest models of inflation which typically produce running of order $\mathcal{O}[\left(1-n_s\right)^2]$. Moreover, a large running also requires further additional parameters in order that inflation does not end too quickly after the observed scales leave the horizon \cite{Easther:2006tv}. The temperature anisotropy excess implied by tensors is also not a smooth function of scale, but rather cut off at the horizon at recombination. To counter this excess, a transition in the scalar power spectrum that occurs more sharply, though coincidentally near these scales, would be preferred. Such a transition can occur without affecting the tensor spectrum if there is a slow-roll violating step in the tensor-scalar ratio while the Hubble rate is left nearly fixed. In this work we consider the effects of placing such a feature near scales associated with the horizon at recombination, thereby suppressing the scalar spectrum on large scales. This slow-roll violating behavior also produces oscillations in the power spectrum \cite{Adams:2001vc, Peiris:2003ff, Park:2012rh, Miranda:2012rm} and generates enhanced non-Gaussianity \cite{Chen:2006xjb, Chen:2008wn} if this transition occurs in much less than an efold. For transitions that alleviate the tensor-scalar tension, these oscillations would violate tight constraints on the acoustic peaks and hence only transitions that occur over at least an efold are allowed. The resulting non-Gaussianity is then undetectable \cite{Adshead:2011bw, Adshead:2012xz}. Throughout, we work in natural units where the reduced Planck mass $M_{\rm Pl} = (8\pi G_N)^{-1/2} = 1$ as well as $c = \hbar = 1$.
A transient violation of slow-roll which generates a step in the scalar power spectrum at scales near to the horizon size at recombination can alleviate problems of predicted excess power in the temperature spectrum, present already in the best fit $\Lambda$CDM spectrum, and greatly exacerbated by tensor contributions implied by the BICEP2 measurement. Such a step may be generated by a sharp change in the speed of the rolling of the inflaton $\epsilon_H$ or by a sharp change in the speed of sound $c_s$ over a period of around an efolding which combine to form the tensor-scalar ratio. Preference for a step from the temperature power spectrum is at a level of $2\Delta\ln {\cal L}_P = -14.2$ if $r=0.2$ and is still $-8.6$ at $r=0.1$, the lowest plausible value that would fit the BICEP2 data. Such an explanation makes several concrete predictions. Since slow-roll is transiently violated in this scenario, there will be an enhancement in the associated three-point correlation function. However, we do not expect this signal to be observable as it impacts only a small number of modes \cite{Adshead:2011bw, Adshead:2012xz}. $E$-mode fluctuations on similar scales would be predicted to have a smaller enhancement then with tensors alone. This prediction should soon be testable; in the BICEP2 data it brings down the total likelihood improvement to $2\Delta\ln {\cal L}_{\rm tot}=-13.7$ with a step at $r=0.2$. While we have used a DBI type Lagrangian to illustrate the impact of a change in the tensor-scalar ratio parameter $\epsilon_H c_s$ due to a step in the sound speed, we do not expect that our results require this form, though precise details of the fit may change. Transient shifts in the speed of sound have been found to occur in inflationary models where additional heavy degrees of freedom have been integrated out \cite{Achucarro:2012yr}. We leave investigation of specific constructions to future work.
14
3
1403.5231
The recent BICEP2 B-mode polarization determination of an inflationary tensor-scalar ratio r=0.2-0.05+0.07 is in tension with simple scale-free models of inflation due to a lack of a corresponding low multipole excess in the temperature power spectrum which places a limit of r<SUB>0.002</SUB>&lt;0.11 (95% C.L.) on such models. Single-field inflationary models that reconcile these two observations, even those where the tilt runs substantially, introduce a scale into the scalar power spectrum. To cancel the tensor excess, and simultaneously remove the excess already present without tensors, ideally the model should introduce this scale as a relatively sharp transition in the tensor-scalar ratio around the horizon at recombination. We consider models which generate such a step in this quantity and find that they can improve the joint fit to the temperature and polarization data by up to 2ΔlnL≈-14 without changing cosmological parameters. Precision E-mode polarization measurements should be able to test this explanation.
false
[ "such models", "models", "cosmological parameters", "simple scale-free models", "tensors", "Precision E-mode polarization measurements", "C.L.", "the temperature power spectrum", "Single-field inflationary models", "the scalar power spectrum", "recombination", "inflation", "the model", "a corresponding low multipole", "r<SUB>0.002</SUB>&lt;0.11", "tension", "an inflationary tensor-scalar ratio", "r=0.2", "the tensor excess", "the temperature and polarization data" ]
11.484388
-0.557014
89
12405086
[ "Bogod, V. M.", "Alissandrakis, C. E.", "Kaltman, T. I.", "Tokhchukova, S. K." ]
2015SoPh..290....7B
[ "RATAN-600 Observations of Small-Scale Structures with High Spectral Resolution" ]
11
[ "St. Petersburg Branch, Special Astrophysical Observatory, Russian Academy of Sciences, 196140, St. Petersburg, Russia; St Petersburg National Research University ITMO, St. Petersburg, Russia; St Petersburg National Research University ITMO, St. Petersburg, Russia", "Section of Astro-Geophysics, Department of Physics, University of Ioannina, 45110, Ioannina, Greece", "St. Petersburg Branch, Special Astrophysical Observatory, Russian Academy of Sciences, 196140, St. Petersburg, Russia", "St. Petersburg Branch, Special Astrophysical Observatory, Russian Academy of Sciences, 196140, St. Petersburg, Russia" ]
[ "2015ARep...59..542P", "2015SoPh..290....1P", "2017A&A...605A..78A", "2018A&A...619L...6N", "2020FrASS...7...74A", "2021FrASS...7...77A", "2022FrASS...9.1320A", "2023Ge&Ae..63.1015I", "2023MNRAS.519.4397S", "2023PhDT........27K", "2023SoPh..298..124M" ]
[ "astronomy" ]
4
[ "Radio emission", "quiet", "Polarization", "radio", "Magnetic fields", "chromosphere", "Chromosphere", "Transition region", "Corona", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1974SoPh...34..125K", "1975IzKry..53..121E", "1975MNRAS.173...65K", "1975SvAL....1...79K", "1975SvAL....1..205B", "1977SvAL....3..300G", "1978ApJ...224.1043Z", "1978SoSAO..23...22B", "1979ApJ...234.1122K", "1982ASSL...96..109G", "1983SoPh...85..237C", "1988ApJ...329..991G", "1990ApJ...355..321G", "1996ApJ...473..539B", "1997A&A...320..993B", "1997ApJ...488..499K", "2004R&QE...47..227B", "2005SoPh..231....1P", "2006A&A...456..697W", "2006A&A...460..901J", "2007ApJ...659.1673A", "2011A&A...532A.112Z", "2011AstBu..66..190B", "2011AstBu..66..205B", "2011AstBu..66..371B", "2011SoPh..273..309S", "2012A&A...537A.152P", "2012ApJ...750L..25J", "2012SSRv..169..181T", "2012SoPh..275....3P", "2012SoPh..275...17L", "2013PASJ...65S..19K" ]
[ "10.1007/s11207-014-0526-6", "10.48550/arXiv.1403.7658" ]
1403
1403.7658_arXiv.txt
\label{intro} The microwave range is rich in information on quiet-sun structures in the upper chromosphere, the transition region and the low corona (see \opencite{2011SoPh..273..309S}, for a recent review). The variation of the brightness temperature with frequency provides valuable data for modeling, while the circular polarization gives estimates of the magnetic field. Due to the limited spatial resolution, useful information can only be obtained with large radio telescopes or interferometers. The chromospheric network was first detected in interferometric data (\opencite{1974SoPh...34..125K}; \opencite{1975MNRAS.173...65K}; \opencite{1978ApJ...224.1043Z}). The first, one-dimensional, imaging observations came from the {\it Radio Astronomical Telescope of the Academy of Sciences} 600 (RATAN-600) in the wavelength range of 2--4 cm, which showed small-scale brightness fluctuations across the quiet-sun (\opencite{1975SvAL....1..205B}; \opencite{1977PAZh....3..550G}; \opencite{1978SoSAO..23...22B}); this phenomenon was called {\it solar radiogranulation}. Thanks to the high spatial resolution (9\arcsec\ by 50\arcsec) at the wavelength of 1.35 cm, it was possible to identify the individual elements of the radiogranulation with bright and dark elements of the chromospheric network visible in the Ca{\sc ii} K line \cite{1982ASSL...96..109G}. In the millimeter range, small-scale structures (SSS) were first observed during the passage of Mercury across the disk of the Sun with the {\it Radio Telescope 22} (RT-22) \cite{1975IzKry..53..121E} and with the {\it Radio Telescope 25} RT-25 by \inlinecite{1975PAZh....1...24K}. \inlinecite{1982ASSL...96..109G} made a comparison with optical images. Typical sizes of the cells obtained by \inlinecite{1978SoSAO..23...22B} and \inlinecite{1982ASSL...96..109G} were 40\arcsec\ to 50\arcsec\ in the range from 1.96 cm to 10 cm. Further observations with two-dimensional resolution came from synthesis instruments. \inlinecite{1979ApJ...234.1122K} were the first to obtain quiet-sun images at the wavelength of 6 cm with arc-second resolution using the {\it Westerbork Synthesis Radio Telescope} (WSRT). The WSRT quiet-sun images showed a clear association of the quiet-sun microwave emission with the chromospheric network. This conclusion was subsequently verified with the {\it Very Large Array} (VLA) observations at 6 and 20\,cm by \inlinecite{1988ApJ...329..991G} and by \inlinecite{1990ApJ...355..321G} at 3.6\,cm. In the mid 90's the VLA was used for quiet-sun observations in the short cm-range (1.2, 2.0 and 3.6\,cm) by \inlinecite{1996ApJ...473..539B}, \inlinecite{1997AA...320..993B}, and \inlinecite{1997ApJ...488..499K}. In addition to the chromospheric network, SSS are also associated with other features, such as X-ray bright points, ephemeral active regions, magnetic elements and, possibly, with small coronal loops. The chromospheric structure is of special interest as an input to the theories of coronal heating (\opencite{2007ApJ...659.1673A}; \opencite{2012A&A...537A.152P}; \opencite{2011A&A...532A.112Z}; see also the review by \opencite{2012SSRv..169..181T}). In recent years significant advances have been made in the study of solar EUV emission with the {\it Hinode} instruments \cite{2012ApJ...750L..25J}, as well as the {\it Atmospheric Imaging Assembly} (AIA) telescopes aboard the {\it Solar Dynamic Observatory} (SDO) \cite{2012SoPh..275...17L}; nevertheless, these instruments provide no information on the strength of the magnetic field. To estimate the magnetic field at the level of the transition region and lower corona, we need to resort to radio observations with the large radio telescope RATAN-600, since the VLA does not have, up to now, sufficient sensitivity to provide accurate measurements of the low circular polarization in the quiet Sun. Despite the lower spatial resolution of RATAN-600, compared to the VLA or to EUV instruments, its sensitivity is high enough to record the SSS in the microwave range. In recent years a new multi-wavelength spectral and polarization equipment has been installed in RATAN-600 (\opencite{2011AstBu..66..190B}; \opencite{2011AstBu..66..205B}), which has a spectral resolution of 1\% and a frequency range from 3\,GHz to 18.2\,GHz. This system uses a high-speed digital data acquisition system with a large dynamic range, thus providing a high signal-to-noise ratio from the level of the quiet Sun up to to that of bursts, in many wavelengths simultaneously. In this article we present the first results of a study of small-scale quiet-sun structures in the microwave range during the time interval from September 2005 to July 2012, made with the RATAN-600 using its new technical capabilities. The observations and their analysis are described in Section 2; our results are presented in Section 3 and are discussed in the Section 4.
In spite of the large width of the RATAN-600 beam in the N--S direction, small scale structures are detected in the 1D scans, in the wavelength range of 1.65 to 10\,cm. The sensitivity of the instrument is sufficient to detect flux variations of $\sim$ 0.01 sfu. The SSS, with a brightness of 1 to 3\% of the quiet-sun level, are stable over several hours, with a characteristic time scale of about one day. Their observed characteristic size increases with wavelength from about 25\arcsec\ to about 50\arcsec, as the resolution decreases; after $\sim7.5$\,cm the characteristic size stabilizes, which could be an indication of a real structural change. A comparison of the microwave SSS with SDO images, treated in a way that simulates the behavior of RATAN-600, showed an almost one-to-one correspondence between the microwave structures and those seen in the 304\,\AA\ AIA band, with a somewhat inferior correlation with the 1600\, \AA\ AIA band; the correlation with the 171\,\AA\ band was small, so we may conclude that the radio SSS emission comes from a region extending from the chromosphere to the low transition region and does not extend up to the height of formation of the 171\,\AA\ band, which has its peak sensitivity near a temperature of $10^6$\,K. The exact identification of the SSS with particular structures proved to be extremely difficult, except in the case of patches of old plage and some intense bipolar regions. The rest of the features appear to be due to the merging of several individual structures of the chromospheric network that exist within the instrumental beam. Our results are in conformity with the conclusion of previous observations (summarized by \opencite{2011SoPh..273..309S}), as far as the association of radio structures with the chromospheric network is concerned, moreover we cover here a more extended frequency range than previous observations. As for the network itself, it is well known from EUV data that it becomes diffuse in the upper transition region and disappears in the low corona; for example, there is no trace of the network in SDO images in the 171\,\AA\ band. A similar behavior is expected for the radio network and our results are consistent with this. In addition to the association of radio SSS with EUV structures, RATAN-600 data provided measurements of the circular polarization and hence of the magnetic field. Our estimates of 40 to 200 Gauss should be considered as a weighted average, due to the integration over the instrumental beam. Although the spatial resolution of our data is much inferior to that obtained by modern instruments operating in the EUV, RATAN-600 observations of SSS have certain merits that we would like to emphasize. The most important one is the possibility to estimate the magnetic field, something that is still beyond the capability of EUV instruments. Moreover, the complete spectral coverage of the RATAN-600 which extends over a frequency range with $f_{max}/f_{min}\sim6$, is an extremely usuful input to detailed modelling of the solar atmosphere from the upper chromosphere to the low corona, where the microwave emission is formed. We note, in this context, that very little use has been made of the microwave fine structure data in multi-component atmospheric models ({\it cf.} \opencite{1983SoPh...85..237C}). Needless is to add that the interpretation of thermal radio emission is devoid of effects such as abundances, excitation and ionization equilibria {\it etc.\/}, that plague the interpretation of EUV data. The new radio facilities currently under development will certainly add to our understanding of the physics of the network and its extent into the upper layers of the solar atmosphere. In particular, accurate polarization measurements will be possible with the {\it Jansky Very Large Array}, while the {\it Atacama Large Millimeter/submillimeter Array} (ALMA) is expected to extend our vision to the sub-mm range ({\it cf.} \opencite{2006AA...456..697W}). The multi-frequency version of the {\it Siberian Solar Radio Telescope} (SSRT), currently under development (see \opencite{2013PASJ...65S..19K}), is expected to provide two-dimensional spectral information. Last but not least, the {\it Expanded Owens Valley Solar Array} will provide a wide frequency coverage, although with a limited number of baselines and, moreover, it will be a solar dedicated instrument. Unfortunately, the ultimate solar radio instrument, the {\it Frequency Agile Solar Radiotelescope} (FASR), is still far from being realized. The observed fine structure of the quiet-sun in microwaves apparently reflects manifestations of plasma inhomogeneities in the form of supergranulation or chromospheric network. This is indicated by their the characteristic dimensions and life times \cite{2005SoPh..231....1P}. As suggested by many authors, the transformation of the magnetic energy into thermal energy may occur in the chromosphere. The main candidate here could be magnetic reconnection, as the most efficient mechanism both for the release of energy into the corona and the transfer of cool plasma from the chromosphere to the corona (\opencite{2012SSRv..169..181T}, \opencite{2011A&A...532A.112Z}). \inlinecite{2007ApJ...659.1673A} considered a number of arguments that point to the heating processes in the chromosphere. One of the main problems here is the magnetic complexity in transition region, so it is very difficult to trace the magnetic connectivity from coronal loops down to the photosphere (\opencite{2006A&A...460..901J}). Future work on fine structures in the chromosphere and low transition region with large optical and ultraviolet telescopes, as well as observations of fine radio structure in microwaves with radio telescopes with big surface area will improve our understanding of their physics and their dynamics. RATAN-600 provides a large data base of such observations, an extended frequency range and the capability of measuring circular polarization. Therefore, it can have an important contribution to their study.
14
3
1403.7658
We present observations of quiet-Sun small-scale structures (SSS) in the microwave range with the Radio Astronomical Telescope of the Academy of Sciences 600 (RATAN-600) spectral-polarization facility in a wide range of frequencies. The SSS are regularly recorded in routine observations of the large reflector-type radio telescope and represent manifestations in the radio range of various structures of the quiet Sun: supergranulation network, bright points, plage patches, and so on. A comparison with images from the Solar Dynamics Observatory (SDO) showed that the microwave emission comes from a region extending from the chromosphere to the low transition region. We measured the properties of the SSS as well as the degree of circular polarization averaged over the beam of the radio telescope, and from this we estimated the magnetic field at the formation level of the radiation.
false
[ "the low transition region", "plage patches", "various structures", "circular polarization", "bright points", "frequencies", "a region", "the radio range", "Sciences", "SSS", "the microwave range", "routine observations", "Sun", "the radio telescope", "a wide range", "manifestations", "the large reflector-type radio telescope", "the Radio Astronomical Telescope", "quiet-Sun small-scale structures", "SDO" ]
12.069888
15.407655
2
437484
[ "Munari, U.", "Mason, E.", "Valisa, P." ]
2014A&A...564A..76M
[ "The narrow and moving HeII lines in nova KT Eridani" ]
12
[ "INAF Osservatorio Astronomico di Padova, 36012, Asiago (VI), Italy", "INAF Osservatorio Astronomico di Trieste, 34143, Trieste, Italy", "ANS Collaboration, c/o Osservatorio Astronomico, 36012, Asiago (VI), Italy" ]
[ "2014A&A...569A..84M", "2015NewA...39...64T", "2016MNRAS.455L..57M", "2016MNRAS.460.2822B", "2017arXiv170909400K", "2018ApJ...857...68H", "2018MNRAS.480..572A", "2020A&ARv..28....3D", "2021gacv.workE..44D", "2022ApJ...941..138W", "2022MNRAS.517.3864S", "2023MNRAS.524.4867I" ]
[ "astronomy" ]
3
[ "novae", "cataclysmic variables", "stars: individual: KT Eridani", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1957gano.book.....G", "1960stat.book..585M", "1975Ap&SS..36..137D", "2008clno.book..232K", "2009ATel.2327....1R", "2009CBET.2055....1R", "2010ATel.2392....1B", "2010ATel.2418....1N", "2010ATel.2423....1B", "2010ATel.2434....1O", "2010ApJ...724..480H", "2011AAS...21733811W", "2011ASPC..451..271H", "2011MNRAS.410..525M", "2011PhDT.........1R", "2012A&A...537A..34J", "2012A&A...544A.149M", "2012BaltA..21...13M", "2012JAVSO..40..582M", "2012PASJ...64..120I", "2013MNRAS.433.1991R", "2013MNRAS.433.2657R", "2014ASPC..490..199M", "2014NewA...27...25M" ]
[ "10.1051/0004-6361/201323180", "10.48550/arXiv.1403.3284" ]
1403
1403.3284_arXiv.txt
Nova Eri 2009, later named KT Eri, was discovered by K. Itagaki on 2009 November 25.5 UT (see CBET 2050), well past its optical maximum. Using data obtained by SMEI (Solar Mass Ejection Imager) on board the {\it Coriolis} satellite, Hounsell et al. (2011) were able to reconstruct the pre-discovery outburst light curve, which highlights a rapid rise in magnitude after the first detection on 2009 November 13.12 UT, a sharp maximum reached on 2009 November 14.67 UT, after which the nova immediately entered the rapid decline characterized by $t_2$=6.6 days. Preliminary reports on the early spectroscopic and photometric evolution were provided by Ragan et al. (2009), Rudy et al. (2009), Bode et al. (2010), Imamura and Tanabe (2012), and Hung, Chen, \& Walter (2012). Radio observations were obtained by O`Brien et al. (2010) and X-ray observations by Bode et al. (2010), Beardmore et al. (2010), and Ness et al. (2010). Raj, Banerjee \& Ashok (2013) discussed early infrared photometric and spectroscopic evolution, while the line profiles and their temporal evolution were modeled in detail by Ribeiro et al. (2013). Jurdana-{\v S}epi{\'c} et al. (2012) searched the Harvard plate archive and measured the progenitor of the nova on 1012 plates dating from 1888 to 1962. No previous outburst was found. The photometric evolution of KT Eri after it returned to quiescence and its persistent P=752 day periodicity have been discussed by Munari and Dallaporta (2014). \begin{figure*} \centering \includegraphics[height=18cm,angle=270]{Figure_1.ps} \caption{Sample spectra from our monitoring to highlight the spectroscopic evolution of KT Eri during the 2009/10 nova outburst and the subsequent return to quiescence.} \label{fig1} \end{figure*} Here we present KT\,Eri spectra taken from outburst maximum to subsequent quiescence and focus on the appearance and evolution of a narrow HeII 4686 \AA\ emission line. Sharp emission lines superimposed to much broader emission components, have been observed in a few other recent novae: YY\,Dor, nova LMC 2009, U\,Sco, DE\,Cir, and V2672\,Oph (see, e.g., the Stony Brooks SMART Atlas\footnote{\sf www.astro.sunysb.edu/fwalter/SMARTS/NovaAtlas/atlas.html}). Complex line profiles have always been modeled with axisymmetric ejecta geometries consisting of bipolar lobes, polar caps, and equatorial rings (starting with Payne-Gaposchkin in 1957). Using a similar approach, the sharp and strong narrow emission in V2672\,Oph could be successfully modeled as coming from an equatorial ring whereas the broader pedestal originates from polar cups (Munari et al. 2011). However, because of their sharpness, profile, and width it has been also suggested that the narrow components in the above systems might arise from the accretion disk of the underlying binary (Walter \& Battisti 2011, but see also Mason \& Walter 2013), once the ejecta becomes sufficiently transparent. In the case of U\,Sco, the observation of radial velocity motion of the narrow HeII emission has been interpreted as restored accretion shortly after the nova 2010 outburst (Mason et al. 2012). Whether a narrow emission component, and in particular the appearance of the sharp HeII$\lambda$4686 line, always originates in the central binary and recovered accretion from the secondary star has to be established. We believe KT Eri offers an interesting bridge between these two alternative views. We will show how when first seen in emission in the spectra of KT Eri, during the optically thick phase, the sharp HeII 4686~\AA\ line was coming from the inner and slower regions of the ejecta, and how, at later times when the ejecta turned optically thin, the HeII line became two times sharper and variable in radial velocity, indicating it was coming directly from the central binary. Thus, the presence and the origin of sharp HeII emission lines seems to depend on the geometry of the ejecta, their viewing angle, and on the evolutionary phase of the nova.
14
3
1403.3284
We present outburst and quiescence spectra of the classical nova KT Eri and discuss the appearance of a sharp HeII 4686 Å emission line, whose origin is a matter of discussion for those novae that showed a similar component. We suggest that the sharp HeII line, when it first appeared toward the end of the outburst optically thick phase, comes from the wrist of the dumbbell structure characterizing the ejecta. When the ejecta turned optically thin, the already sharp HeII line became two times narrower and originated from the exposed central binary. During the optically thin phase, the HeII line displayed a large change in radial velocity that had no counterpart in the Balmer lines (both their narrow cores and the broad pedestals). The large variability in radial velocity of the HeII line continued well into quiescence, and it remains the strongest emission line observed over the whole optical range.
false
[ "HeII", "the sharp HeII line", "radial velocity", "the HeII line", "a sharp HeII 4686 Å emission line", "the strongest emission line", "KT Eri", "the already sharp HeII line", "the Balmer lines", "the exposed central binary", "quiescence", "the whole optical range", "HeII 4686 Å", "discussion", "Balmer", "a similar component", "the classical nova KT Eri", "nova", "the broad pedestals", "the dumbbell structure" ]
5.143344
9.38196
25
3169222
[ "Krumholz, Mark R." ]
2015ASSL..412...43K
[ "The Formation of Very Massive Stars" ]
38
[ "Department of Astronomy &amp; Astrophysics, University of California" ]
[ "2014A&A...568L..13W", "2015A&A...580A..92R", "2015ApJ...805...92O", "2015ApJ...806L..34F", "2015MNRAS.451.3034N", "2016A&A...587A.115M", "2016A&A...588L...7K", "2016A&A...592A..77B", "2016ASPC..507..147D", "2016ApJ...824...94R", "2017A&A...602A..77T", "2017A&A...607A..22R", "2017ApJ...846..100G", "2017NatAs...1..784O", "2017hsn..book..567H", "2018A&A...615A.119V", "2018ARA&A..56...41M", "2018MNRAS.477.5422A", "2018MNRAS.478.5037R", "2018MNRAS.480.3916M", "2018sf2a.conf..311L", "2019A&A...625L...2K", "2019FrASS...6...54P", "2020A&A...635A.163B", "2020A&A...642A..87K", "2020MNRAS.496.4311C", "2020MNRAS.499.1918B", "2020pase.conf..105A", "2021A&A...648A..66G", "2021ApJ...921...21H", "2021ExA....51..887G", "2021MNRAS.501.5257R", "2021PASJ...73S.405F", "2022A&A...659A..81B", "2022A&A...662A...8M", "2022ApJ...925....3S", "2023A&A...674A.160G", "2023ApJ...944..211L" ]
[ "astronomy" ]
29
[ "Physics", "Astrophysics - Solar and Stellar Astrophysics", "Astrophysics - Astrophysics of Galaxies" ]
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[ "10.1007/978-3-319-09596-7_3", "10.48550/arXiv.1403.3417" ]
1403
1403.3417_arXiv.txt
\label{sec:intro} The mechanism by which the most massive stars form, and whether there is an upper limit to the mass of star that this mechanism can produce, has been a problem in astrophysics since the pioneering works of \citet{larson71a} and \citet{kahn74a}. These authors focused on the physical mechanisms that might inhibit accretion onto stars as they accreted interstellar matter, and we will return to this topic below. However, a more modern approach to the problem of very massive stars requires placing them in the context of a broader theory of the stellar initial mass function (IMF). The IMF is characterized by a peak in the range $0.1-1$ $\msun$, and a powerlaw tail at higher masses of the form $dn/d\ln m\propto m^\Gamma$ with $\Gamma\approx -1.35$ \citep[and references therein]{bastian10a}. However, the mass to which this simple powerlaw extends is not very well-determined. It is not possible to measure the IMF for field stars to very high masses due to uncertainties in star formation histories and the limited number of very massive stars available in the field. Measurements of the high-mass end of the IMF in young clusters must target very massive systems in order to achieve strong statistical significance, and such clusters are rare and thus distant. This creates significant problems with confusion. The limited studies that are available suggest that the a powerlaw with $\Gamma \approx -1.35$ remains a reasonable description of the IMF out to masses of $\sim 100$ $\msun$ or more \citep[e.g.][]{massey98b, kim06b, espinoza09a}. However, it is by no means implausible that there are hidden features lurking in the IMF at the highest masses. Indeed some analyses of the IMF have claimed to detect an upper cutoff (see the Chapter by F.~Martins in this volume for a critical review). This observational question of whether the most massive stars are simply the ``tip of the iceberg" of the normal IMF, or whether they represent a fundamentally distinct population, animates the theoretical question about how such stars form. The two dominant models for how massive stars form are formation by accretion of interstellar material, i.e.~the same mechanism by which stars of low mass form, and formation by collisions between lower mass stars, which would represent a very different formation mechanism from the bulk of the stellar population.\footnote{Mergers between two members of a tight binary as a result of the growth of stellar radii during main sequence or post-main sequence evolution, or as a result of secular interactions in hierarchical triples, is a third possible mechanism by which massive stars can and probably do gain mass \citep{sana12a, de-mink13a, moeckel13a}. However, I do not discuss this possibility further, because it provides at most a factor of two increase in stellar mass.} In the remainder of this Chapter, I review each of these models in turn (Sections \ref{sec:accretion} and \ref{sec:collision}), pointing out its strengths, weaknesses, and areas of incompleteness. I then discuss the observable predictions made by each of these models, and which might be used to discriminate between them (Section \ref{sec:discrimination}). Finally, I summarize and return to the question first raised by \citet{larson71a} and \citet{kahn74a}: is there an upper mass limit for star formation, and if so, why (Section \ref{sec:masslimit})?
\label{sec:masslimit} Having discussed the two main formation scenarios, I now return to the question of whether star formation has a mass limit. To review, there is at present no really convincing evidence that any mechanism is capable of halting the growth of stars by accretion. The classical mechanism for limiting stellar masses is radiation pressure, but non-spherical accretion, produced by some combination of disks, outflow cavities, and instabilities appears to defeat this limit. Similarly, the problem of gas fragmenting too strongly to form massive stars appears to be solved by a combination of radiative heating and magnetic fields, though the possibility that disk fragmentation might at some point limit stellar masses remains. Photoionization and stellar winds are somewhat more promising as mechanisms that might limit stars' growth, but these remain at best possibilities. There are no real analytic models applicable to present-day (as opposed to primordial) star formation that suggest what limits these mechanisms might impose, and there are no simulations demonstrating that either of these processes are capable of terminating accretion. A fair description of the state of the field a decade ago might have been that the presumption was in favor of feedback limiting massive star formation, and that the burden of proof was on those trying to show that feedback could be overcome. The last decade of work has reversed that situation, with all tests thus-far performed showing that accretion is very difficult to reverse. This does not prove that no mechanism can limit stellar masses, but does mean that such a limit would need to be demonstrated. For collisions, the question is not whether but where they can create very massive stars. There is no doubt that collisions and collisional growth will occur if the conditions are dense enough, and the only question is the frequency with which such dense conditions are created in nature, which in turn will determine the contribution of the collisional formation channel to the overall population of very massive stars. No presently-observed star cluster has a density high enough for collisions to be likely, but it is possible that these clusters experienced a very dense phase during which collisions occurred. This could have been either an early gas-dominated phase or a later phase of core collapse aided by primordial mass segregation and high levels of primordial substructure. However, the threshold density required to achieve significant collisional growth depends on details of massive star mergers and winds that are poorly understood. Even for favorable assumptions about these uncertain parameters, it is not clear that the observed present-day properties of massive star clusters can be reconciled with an evolutionary history in which they were once dense enough to have produced collisions. There are a number of observational tests that may be able to settle the question of which of these mechanisms is the dominant route to the formation of the most massive stars. Accretion models predict that massive stars are simply the tip of the iceberg of normal star formation, so that the high end of the stellar mass function is continuous and does not depend radically on the environment, and that massive stars are likely to have low-mass as well as high-mass companions. Although the observable consequences of the collisional formation models have received somewhat less attention, such models appear to predict quite different results: there should be a large gap in stellar mass functions separating the bulk of the accretion-formed stellar population from the few collisionally-formed stars, and this feature should appear only in the most massive and densest clusters. It seems likely that these collisionally-formed stars will lack low-mass companions. It should be possible to perform most or all of these observational tests with the coming generation of telescopes and instruments, which will provide higher angular resolution and contrast sensitivity than have previously been possible. \begin{acknowledgement} I thank all the authors who provided figures for this review: H.~Baumgardt, J.~Dale, N.~Moeckel, A.~C.~Myers, and D.~Vanbeveren. During the writing of this review I was supported by NSF CAREER grant AST-0955300, NASA ATP grant NNX13AB84G, and NASA TCAN grant NNX14AB52G. I also thank the Aspen Center for Physics, which is supported by NSF Grant PHY-1066293, for hospitality during the writing of this review. \end{acknowledgement}
14
3
1403.3417
In this chapter I review theoretical models for the formation of very massive stars. After a brief overview of some relevant observations, I spend the bulk of the chapter describing two possible routes to the formation of very massive stars: formation via gas accretion, and formation via collisions between smaller stars. For direct accretion, I discuss the problems of how interstellar gas may be prevented from fragmenting so that it is available for incorporation into a single very massive star, and I discuss the problems presented for massive star formation by feedback in the form of radiation pressure, photoionization, and stellar winds. For collision, I discuss several mechanisms by which stars might be induced to collide, and I discuss what sorts of environments are required to enable each of these mechanisms to function. I then compare the direct accretion and collision scenarios, and discuss possible observational signatures that could be used to distinguish between them. Finally, I come to the question of whether the process of star formation sets any upper limits on the masses of stars that can form.
false
[ "massive star formation", "star formation", "stars", "smaller stars", "formation", "stellar winds", "gas accretion", "very massive stars", "possible observational signatures", "several mechanisms", "radiation pressure", "a single very massive star", "collision scenarios", "direct accretion", "photoionization", "environments", "incorporation", "collision", "collisions", "feedback" ]
10.778033
10.088654
-1
527222
[ "van den Bosch, Frank C.", "Jiang, Fangzhou" ]
2014arXiv1403.6835V
[ "Statistics of Dark Matter Substructure: II. Comparison of Model with Simulation Results" ]
16
[ "-", "-" ]
[ "2014MNRAS.445.1713V", "2015A&A...575A..59S", "2015ApJ...810...21M", "2015ApJ...812....9E", "2015JETP..121.1104C", "2015MNRAS.447.3693K", "2015MNRAS.452..444C", "2015MNRAS.453.3575J", "2016A&A...587A..24S", "2016A&A...595A..56S", "2016ApJ...820...85Z", "2016MNRAS.455..158V", "2016MNRAS.457.1208H", "2016PhDT.......338M", "2017JCAP...04..018H", "2022JCAP...10..021L" ]
[ "astronomy" ]
7
[ "Astrophysics - Cosmology and Nongalactic Astrophysics", "Astrophysics - Astrophysics of Galaxies" ]
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[ "10.48550/arXiv.1403.6835" ]
1403
1403.6835_arXiv.txt
\label{Sec:Introduction} Hierarchical structure formation in a $\Lambda$CDM cosmology gives rise to dark matter haloes with abundant substructure in the form of self-bound clumps of matter. These subhaloes are the remnants of dark matter haloes that have been accreted by their host halo over cosmic time, and that have (thus far) survived tidal destruction. Dark matter subhaloes host satellite galaxies, boost the dark matter annihilation signal, cause tidal heating of stellar streams and disks, and are responsible for time-delays and flux-ratio anomalies in gravitational lensing. Hence, characterizing the abundance, spatial distribution and internal structure of dark matter substructure is important for a large number of astrophysical applications. Given the highly non-linear nature of hierarchical structure formation, the substructure of dark matter haloes is best studied using high-resolution $N$-body simulations. Ever since the first numerical simulations reached sufficient resolution to resolve dark matter subhaloes (Tormen 1997; Ghigna \etal 1998; Tormen, Diaferio \& Syer 1998; Klypin \etal 1999; Moore \etal 1999) numerous studies have used $N$-body simulations of ever increasing size and/or numerical resolution to study their statistics as function of host halo mass, redshift, cosmology, and other properties of relevance, such as the formation time of the halo (e.g., Springel \etal 2001, 2008; Stoehr \etal 2002; Gao \etal 2004; De Lucia \etal 2004; Diemand, Moore \& Stadel 2004; Gill \etal 2004a,b; Kravtsov \etal 2004; Reed \etal 2005; Shaw \etal 2007; Giocoli \etal 2008a, 2010; Weinberg \etal 2008; Angulo \etal 2009; Boylan-Kolchin \etal 2010; Klypin, Trujillo-Gomez \& Primack 2011; Wu \etal 2013). These studies not only used different simulation codes, different cosmologies, different numerical resolutions, and different simulation volumes, but also different subhalo finders. To date, more than a dozen different subhalo finders have been used, all based on some of the following two characteristics of a subhalo: (i) it is a self-bound, overdense region inside its host halo, and (ii) it was it's own host halo before it merged into its current host (see Han \etal 2012). Most halo finders only use the instantaneous particle positions and velocities to identify subhaloes based on the first characteristic listed above. Most of these only use the velocity information to remove unbound particles from a list of candidate particles identified based on density alone. Examples of such halo finders are \Subfind (Springel \etal 2001), {\tt SKID} (Stadel 2001), Bound Density Maximum (\BDM; Klypin \& Holtzman 1997), Amiga Halo Finder ({\tt AHF}; Knollmann \& Knebe 2009), and Voronoi Bound Zones ({\tt VOBOZ}; Neyrinck, Gnedin \& Hamilton 2005). Others, such as 6-Dimensional Friends-of-Friends ({\tt 6DFOF}; Diemand Kuhlen \& Madau 2006), Hierarchical Structure Finder ({\tt HSF}; Maciejewski \etal 2009) and \Rockstar (Behroozi \etal 2013a), identify (sub)haloes using the full six-dimensional phase-space information. Finally, there are also subhalo finders that make additional use of the second characteristic listed above by employing the time domain. Since subhaloes are remnants of dark matter host haloes, one can identify the former by tracing the member particles of the latter that remain part of a self-bound entity. Examples of these are \SURV (Tormen \etal 1998) and the Hierarchical Bound-Tracing algorithm ({\tt HBT}) of Han \etal (2012). Note that \Rockstar also uses some time-domain information in its (sub)halo identification, making it the only subhalo finder that uses information in seven dimensions. In an era of precision cosmology, in which accurate, percent level characterization of the abundances of dark matter haloes and subhaloes is crucial, it is of paramount importance to compare the performance of all these different subhalo finders, and to quantify their accuracy and reliability. Unfortunately, and somewhat surprisingly, this has received relatively little attention. Muldrew, Pearce \& Power (2011) compared the performances of \Subfind and {\tt AHF} in recovering mock subhaloes placed in a mock host halo. They showed that the mass of the subhalo recovered by \Subfind has a strong dependence on its radial position within the host halo, and that neither subhalo finder can accurately recover the subhalo mass when it is near the center of the host halo. More quantitatively, \Subfind was only able to recover 50 percent of the subhalo mass when its center is located at half the virial radius from the center of its host. At $r < r_{\rm vir}/10$, neither subhalo finder could recover more than 40 percent of the actual subhalo mass. These problems arise from the subhalo being defined as a mere overdensity in configuration space. Indeed, using a similar approach based on mock haloes, Knebe \etal (2011) showed that this problem can be significantly reduced (but not eliminated) when using a subhalo finder that operates in 6D phase-space. A potential problem with these two studies, though, is that they used mock haloes. As pointed out in Knebe \etal (2011), the discrepancy between the true and recovered subhalo masses is likely to be overestimated in this idealized setup. In reality, a subhalo experiences tidal stripping and truncation when moving towards the center of its host halo, and the mass discrepancy is likely to be reduced when only considering the mass within the tidal truncation radius. Following up on the initial study by Knebe \etal (2011), Onions \etal (2012) therefore compared the performance of subhalo finders using an ultra-high resolution simulation of a single Milky-Way sized dark matter halo from the Aquarius project (Springel \etal 2008). Comparing the statistics and properties of the dark matter subhaloes identified within this single host halo with ten different subhalo finders, and using a common post-processing pipeline to uniformly analyze the particle lists provided by each finder, they find that the basic properties (mass and maximum circular velocity) of dark matter subhaloes can be reliably recovered (at an accuracy better than 20 percent) if the subhaloes contain more than 100 particles. In a follow-up study, Knebe \etal (2013) showed that discarding the results from the two subhalo finders that lack a (reliable) procedure to remove unbound particles, the scatter among the (eight remaining) subhalo finders is reduced by a factor two, to $\sim 10\%$. Finally, the studies of Onions \etal (2012) and Knebe \etal (2013) show that configuration finders yield less reliable masses for subhaloes close to the center of their host than phase-space finders, but that the differences are significantly smaller than suggested by the tests based on mock haloes described above. Unfortunately, since the study by Onions et al. only used a single dark matter halo, albeit at exquisite numerical resolution, the comparison is limited to the relatively low mass end of the subhalo mass function, where the cumulative mass function $N(>m)$, exceeds unity. In order to study the massive end of the subhalo mass function, where $N(>m) < 1$, one needs to average over large numbers of host haloes. The abundances of these rare but massive subhaloes has important implications for, among others, the statistics of massive satellite galaxies (e.g., Boylan-Kolchin \etal 2010; Busha \etal 2011) and the detection rate of dark matter substructure via gravitational lensing (e.g., Vegetti \etal 2010, 2012). In this paper we use subhalo mass functions and subhalo catalogs from a variety of numerical simulations that are publicly available, and that have been obtained using different subhalo finders, to compare subhalo mass functions, focusing on the massive end. We confirm the findings by Onions et al., that the subhalo mass functions are consistent at the 20 percent level at the low-mass end. At the massive end, though, different subhalo finders yield subhalo abundances that differ by more than one order of magnitude! By comparing the simulation results with a new, semi-analytical model (Jiang \& van den Bosch 2014b), we demonstrate that subhalo finders that identify subhaloes based purely on density in configuration space, such as the popular \Subfind and \BDM, dramatically underpredict the masses, but not the maximum circular velocities, of massive subhaloes. We also show that the model predictions are in excellent agreement with the simulation results when they are analyzed using more advanced subhalo finders that use phase-space and/or time domain information in the identification of subhaloes. We discuss a number of implications of our findings, in particular with regard to the power-law slope of the subhalo mass and velocity functions. Throughout we use $m$ and $M$ to refer to the masses of subhaloes and host haloes, repectively, use $\ln$ and $\log$ to indicate the natural logarithm and 10-based logarithm, respectively, and express units that depend on the Hubble constant in terms of $h = H_0/(100\kmsmpc)$.
\label{sec:conc} We have compared subhalo mass and velocity functions obtained from different simulations, with different subhalo finders, among each other and with predictions from our new semi-analytical model presented in Paper~I. Our findings can be summarized as follows: \begin{itemize} \item We confirm the findings of Onions \etal (2012) and Knebe \etal (2013) that the subhalo mass functions obtained using different subhalo finders agree with each other at the level of $\sim 20$ percent. However, this is only true for low-mass subhaloes with $m/M \lta 0.1$; at the more massive end, different subhalo finders yield SHMFs that differ by more than an order of magnitude! \item Subhalo finders that identify subhaloes based purely on density in configuration space, such as \Subfind and \BDM, dramatically underpredict, by more than an order of magnitude, the abundances of massive subhaloes (with masses $m \gta M/10$), especially in more massive host haloes. These problems are much less severe for the subhalo velocity functions, indicating that they arise from issues related to assigning masses to the subhaloes, rather than from issues related to detecting them. The maximum circular velocity probes the inner regions of dark matter haloes, and can therefore be measured much more reliably than subhalo mass. \item Overall our model predictions are in excellent agreement with simulation results obtained using the more advanced subhalo finders \Rockstar and \SURV. In particular, the model accurately reproduces the slope and host-mass-dependent normalization of both the subhalo mass and velocity functions. There are small discrepancies at the very massive end, but rather than reflecting an inaccuracy of the model, these arise from subtle issues having to do with the exact halo mass definitions of overlapping haloes. \item Since tidal stripping and heating impact the outskirts of subhaloes much more than their inner regions, a large reduction in mass only has a relatively mild impact on the maximum circular velocity (Hayashi \etal 2003; Pe\~narrubia \etal 2008, 2010). We confirm our findings from Paper~I that, on average, the relation between $V_{\rm max}/V_{\rm acc}$ and $m/m_{\rm acc}$ is well described by Eq.~(\ref{Vmaxfit}) with $(\eta,\mu) = (0.44,0.60)$, which is roughly bracketed by the relations obtained by Hayashi \etal (2003) and Pe\~narrubia \etal (2010) using idealized $N$-body simulations. However, there are small but noticeable differences in the best-fit values of $\eta$ and $\mu$ for different subhalo finders, indicating that not only the abundances, but also the structural properties of dark matter subhaloes, are sensitive to the subhalo finder used. \item The mass and velocity functions obtained from the Bolshoi and MultiDark simulations confirm our finding from Paper~I that the power-law slopes of $\rmd N/\rmd\log(m/M)$ and $\rmd N/\rmd\log(V_{\rm max}/V_{\rm vir})$ are with $0.82$ and $2.6$, respectively, significantly shallower than what has been claimed in several studies in the literature. In particular, studies based on \Subfind by Boylan-Kolchin \etal (2010) and Gao \etal (2012) have yielded slopes that are significantly steeper. Given the excellent agreement between our model predictions and the \Rockstar, \BDM and \SURV results, we believe that this discrepancy reflects a problem with \Subfind. We emphasize that accurate knowledge of the power-law slope of the subhalo mass and velocity functions is important for calculating the `boost' factor for dark matter annihilation due to substructure, as this requires extrapolation of $\rmd N/\rmd\log(m/M)$ down to the dark matter Jeans mass. \item Comparing the velocity functions for subhaloes of different order, we find that our model is in excellent agreement with the Bolshoi and MultiDark results for first-order subhaloes. This agreement, however, rapidly deteriorates with increasing order; not only between model and simulations, but also among the simulation results themselves. We speculate that these discrepancies are mainly a manifestation of subtle issues having to do with how different subhalo finders treat overlap among haloes and subhaloes. More detailed studies are required to investigate these issues further, and to provide a more reliable testbed for our model predictions. \end{itemize}
14
3
1403.6835
We compare subhalo mass and velocity functions obtained from different simulations with different subhalo finders among each other, and with predictions from the new semi-analytical model of Jiang &amp; van den Bosch (2014). We find that subhalo mass functions (SHMFs) obtained using different subhalo finders agree with each other at the level of ~ 20 percent, but only at the low mass end. At the massive end, subhalo finders that identify subhaloes based purely on density in configuration space dramatically underpredict the subhalo abundances by more than an order of magnitude. These problems are much less severe for subhalo velocity functions (SHVFs), indicating that they arise from issues related to assigning masses to the subhaloes, rather than from detecting them. Overall the predictions from the semi-analytical model are in excellent agreement with simulation results obtained using the more advanced subhalo finders that use information in six dimensional phase-space. In particular, the model accurately reproduces the slope and host-mass-dependent normalization of both the subhalo mass and velocity functions. We find that the SHMFs and SHVFs have power-law slopes of 0.82 and 2.6, respectively, significantly shallower than what has been claimed in several studies in the literature.
false
[ "subhalo mass functions", "different subhalo finders", "subhalo velocity functions", "subhalo finders", "configuration space", "different simulations", "magnitude", "subhalo mass and velocity functions", "simulation results", "several studies", "Bosch", "van den", "excellent agreement", "the new semi-analytical model", "the low mass end", "masses", "the subhalo abundances", "the semi-analytical model", "SHVFs", "density" ]
10.993673
5.16578
131
746403
[ "Raymond, S. N.", "Bonsor, A." ]
2014MNRAS.442L..18R
[ "Vega's hot dust from icy planetesimals scattered inwards by an outward-migrating planetary system." ]
20
[ "CNRS, Laboratoire d'Astrophysique de Bordeaux, UMR 5804, F-33270 Floirac, France; Laboratoire d'Astrophysique de Bordeaux, Univ. Bordeaux, UMR 5804, F-33270 Floirac, France; NASA Astrobiology Institute's Virtual Planetary Laboratory, Seattle, WA 98195, USA", "H.H. Wills Physics Laboratory, School of Physics, University of Bristol, Tyndall Avenue, Bristol BS8 1TL, UK" ]
[ "2014A&A...570A.128E", "2014ApJ...796..141B", "2014MNRAS.441.2380B", "2016ApJ...816...50R", "2016P&SS..133...47M", "2017MNRAS.464.3385W", "2017MNRAS.465.2352F", "2018AJ....156..214M", "2018exha.book.....P", "2019A&A...626A...2S", "2019MNRAS.487.2191D", "2020ApJ...898..146M", "2020MNRAS.498.2798P", "2021AJ....161..157H", "2021ApJ...923...90M", "2021MNRAS.503.4767P", "2022MNRAS.514.3815M", "2022MNRAS.517.1436P", "2023A&A...678A.121S", "2024ApJ...962L..21K" ]
[ "astronomy" ]
6
[ "planets and satellites: detection", "planets and satellites: dynamical evolution and stability", "zodiacal dust", "stars: Individual: Vega (alpha lyrae)", "planetary systems", "infrared: planetary systems", "Astrophysics - Earth and Planetary Astrophysics", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1969JGR....74.2531D", "1984ApJ...278L..23A", "1994Icar..108...18L", "1998Natur.392..788H", "1999MNRAS.304..793C", "2002ApJ...569L.115W", "2003ApJ...598..626D", "2004Icar..170..492G", "2005ApJ...628..487S", "2006A&A...452..237A", "2006ApJ...641..556M", "2006ApJ...652.1729I", "2007ApJ...663..365W", "2008ARA&A..46..339W", "2008ApJ...688..583H", "2008SPIE.7014E..18B", "2008SPIE.7015E..18M", "2009Icar..199..197K", "2009MNRAS.399..385B", "2010A&A...518L.130S", "2010ApJ...708...71Y", "2010ApJ...708.1728M", "2010ApJ...711..772R", "2010ApJ...713..816N", "2010RAA....10..383K", "2011A&A...530A..62R", "2011A&A...534A...5D", "2012A&A...541A..11R", "2012A&A...548A.104B", "2012ApJ...758...80O", "2013A&A...555A.146L", "2013ApJ...763..118S", "2013MNRAS.429.2894R", "2013MNRAS.433.2938B", "2014MNRAS.441.2380B" ]
[ "10.1093/mnrasl/slu048", "10.48550/arXiv.1403.6821" ]
1403
1403.6821_arXiv.txt
Vega is the original rock star. Observations dating back 30 years have shown a large excess in flux over the stellar photosphere at mid-infrared and longer wavelengths~\citep{aumann84,wilner02,sibthorpe10}. Additional excesses have been measured in the $K$ band ($\sim 2 \micron$) of $\sim 1.3\%$ the stellar flux~\citep{absil06,defrere11} and at $10-30 \micron$ (excess of $\sim 7\%$)~\citep{su13}. The bright emission at $\lambda \ge 50 \micron$ comes from a resolved disk of more than 100~AU in radius~\citep{holland98,su05}. Cold dust is thought to be generated by the collisional grinding of a planetesimal belt located between roughly 60 and 120~AU~\citep{su05,muller10}. The excess at $2 \micron$ is attributed to hot ``exozodiacal'' dust at $\sim 0.2$~AU that is dominated by small grains and cannot be in collisional equilibrium because of the too-short timescales~\citep{defrere11}. The warm excess at $10-30 \micron$ is spatially distinct from the hot emission and may represent an asteroid belt at $\sim 14$~AU~\citep{su13}, although the flux is likely contaminated by the tail of the hot dust~\citep[e.g.][]{lebreton13} and, as we will explore below, perhaps by inward-scattered cold planetesimals. Here we propose that Vega's hot exozodiacal dust may be supplied from its cold debris disk. Planetesimals are transported inward by gravitational scattering by a system of planets. The outermost planet scatters planetesimals inward and passes them off to the next planet until they reach the inner system; this is the same mechanism that delivers Jupiter-family comets to the inner Solar System~\citep{levison94} and is thought to be the main contributor to the Solar System's zodiacal dust~\citep{nesvorny10,rowanrobinson13}. Planets on static orbits cannot sustain the inward-scattering rates needed to produce the observed exozodis because the outer planetesimal supply zones are depleted too rapidly~\citep{bonsor12}. The observed exozodis can also not be caused by dynamical instabilities among systems of giant planets~\citep[see][]{raymond10} because the pulse of inward scattering is too short in duration~\citep{bonsor13}. For hot inner dust to come from the outer planetesimal belt, planets must continually interact with new material. A solution is for the outermost planet to migrate outward. For a given range of planet masses and disk properties, outward migration is naturally maintained by the gravitational back-reaction of inward-scattering of planetesimals~\citep{gomes04,kirsh09,ormel12}. Planets that are too low-mass migrate rapidly but do not excite planetesimal eccentricities sufficiently to produce sufficient scattering. Planetesimal driven migration naturally produces exozodis for a specific range of planet masses that depend on the disk's surface density~\cite[see][]{bonsor14}. Our simulations constitute a proof-of-concept experiment. We determine the conditions needed for planetesimal-driven migration of a system of planets to generate a sufficient flux of inward-scattered planetesimals to roughly reproduce Vega's exozodi. We consider the outer debris disk as the sole source of planetesimals. We show that inward-scattered planetesimals can produce sufficient warm dust to explain the $10-30\micron$ excess without invoking steady-state collisional grinding in a dynamically cold asteroid belt.
We have shown that Vega's exozodi can be explained by a system of planets with orbital radii of roughly 5-60 AU. Planets scatter planetesimals from the outer planetesimal disk into the inner planetary system. The outermost planets undergo plantesimal-driven migration outward into the disk, continually encounter fresh material and can sustain a substantial inward-scattering rate for the age of the system and beyond~\citep[see][for details of this process]{bonsor14}. For the mechanism to operate the system must meet the following requirements. First, the outermost 1-2 planets must have masses such that planetesimal-driven migration is triggered and sustained. Successful simulations had $M_{outer} = 2.5-20 \mearth$. The inner planets' masses have only a minor effect, although the highest scattering rates were for simulations with higher-mass (Saturn- to Jupiter-mass) planets closer-in. Second, while the inner system may host a gas giant, a $\sim$Jupiter-mass planet cannot orbit beyond $\sim15$~AU as it too-efficiently ejects planetesimals from the system and breaks the inward-scattering chain (see Fig.~\ref{fig:jup}). This prediction does not violate current constraints~\citep{itoh06,marois06,heinze08} and may be tested with upcoming instruments. Third, our model requires that a substantial fraction of the mass in planetesimals scattered within 3~AU must end up as fine-grained dust. This fraction is $\sim$10-30\% for the simulations we have deemed successful. If we take our simplified dust flux calculation at face value, our simulations provide a modestly good match to the observed SED (see Fig.~\ref{fig:SED}). Although our simulations somewhat overestimate the flux at $25-50 \micron$, dust produced by in-scattered planetesimals may offer an alternate explanation for the $10-30\micron$ excess detected by \cite{su13}. Planetesimal-driven migration of a system of 5-7 planets in a lower-mass disk and colder disk can also yield reasonable inward-scattering rates and SEDs. \vskip .1in \noindent We thank M. Booth and J.-C. Augereau for interesting discussions. A.B. acknowledges the support of the ANR-2010 BLAN-0505- 01 (EXOZODI) and funding from NERC. This work was performed as part of the NASA Astrobiology Institute's Virtual Planetary Laboratory Lead Team, supported by the NASA under Cooperative Agreement No. NNA13AA93A. This paper was entirely written while S.N.R. was standing and is therefore a ``sit-free'' paper.
14
3
1403.6821
Vega has been shown to host multiple dust populations, including both hot exozodiacal dust at sub-au radii and a cold debris disc extending beyond 100 au. We use dynamical simulations to show how Vega's hot dust can be created by long-range gravitational scattering of planetesimals from its cold outer regions. Planetesimals are scattered progressively inwards by a system of 5-7 planets from 30 to 60 au to very close-in. In successful simulations, the outermost planets are typically Neptune mass. The back-reaction of planetesimal scattering causes these planets to migrate outwards and continually interact with fresh planetesimals, replenishing the source of scattered bodies. The most favourable cases for producing Vega's exozodi have negative radial mass gradients, with sub-Saturn- to Jupiter-mass inner planets at 5-10 au and outer planets of 2.5 - 20 M<SUB>⊕ </SUB>. The mechanism fails if a Jupiter-sized planet exists beyond ∼15 au because the planet preferentially ejects planetesimals before they can reach the inner system. Direct-imaging planet searches can therefore directly test this mechanism.
false
[ "∼15 au", "planetesimal scattering", "fresh planetesimals", "Jupiter-mass inner planets", "negative radial mass gradients", "Planetesimals", "planetesimals", "Neptune mass", "scattered bodies", "Direct-imaging planet searches", "multiple dust populations", "Jupiter", "the outermost planets", "sub", "the planet", "these planets", "a Jupiter-sized planet", "its cold outer regions", "-", "100 au" ]
9.059913
13.286751
-1
678336
[ "Anderson, Gemma J.", "Regan, Donough", "Seery, David" ]
2014JCAP...07..017A
[ "Optimal bispectrum constraints on single-field models of inflation" ]
5
[ "Astronomy Centre, University of Sussex, Falmer Campus, Brighton, BN1 9QH U.K.", "Astronomy Centre, University of Sussex, Falmer Campus, Brighton, BN1 9QH U.K.", "Astronomy Centre, University of Sussex, Falmer Campus, Brighton, BN1 9QH U.K." ]
[ "2013PhDT.......145S", "2014JCAP...10..005G", "2015JCAP...02..015R", "2015MNRAS.448.2232R", "2016PhRvD..94d3510B" ]
[ "astronomy" ]
4
[ "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
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[ "10.1088/1475-7516/2014/07/017", "10.48550/arXiv.1403.3403" ]
1403
1403.3403_arXiv.txt
Successive microwave-background surveys have accumulated some evidence for the inflationary paradigm, in which structure in the universe was seeded by quantum fluctuations during an epoch preceding the hot, dense phase where nucleosynthesis occurred~\cite{Ade:2013ydc,Hinshaw:2012aka}. But despite broad support for the overall framework, attempts to identify the precise degrees of freedom whose quantum fluctuations were relevant have met with less success. Whatever microphysics underlay the putative inflationary epoch remains mysterious. In scattering experiments, an abundance of observables---including, among others, branching ratios, decay rates, and differential dependence on energy or angles---% allow indirect access to microphysical information through reconstruction of the correlation functions, or `$n$-point functions'. These measure interference between quantum fluctuations and encode information about the dynamics of the theory. It is the rich information which can be obtained from reconstruction of the correlation functions which makes measurements in particle physics so constraining. In cosmology our observables are more limited and so is the degree to which the $n$-point functions can be reconstructed. Over a narrow range of scales, the $n$-point functions of the cosmic microwave background (`CMB') anisotropies are sensitive to the $n$-point functions of the primordial `curvature perturbation', which is a calculable, model-dependent mix of the fluctuations imprinted on the light fields of the inflationary epoch. This correspondence has been used for many years to place restrictions on the inflationary model space from measurements of the CMB temperature and polarization two-point functions. But if a \emph{three}-point function of the CMB anisotropies could be measured it would provide access to more nuanced and discriminating microphysical information. Ideally we would like to observe systematic relationships between the $n$-point functions which would point clearly to a quantum mechanical origin for the fluctuations. This is important because it is unclear whether we could ever rule out a non-quantum origin (perhaps associated with new but non-inflationary physics at early times) using only the two-point function. Measurements of the CMB temperature anisotropy have now reached sufficient accuracy that it is feasible to estimate the three-point temperature autocorrelation function. The most precise constraints come from the Planck2013 dataset~\cite{Ade:2013ydc}. But despite the quality of the measurements, the signal-to-noise for any particular combination of wavenumbers is still too low to allow the three-point function to be reconstructed directly. Instead, measurements are made by picking an Ansatz or `template' for the way in which the correlations change with wavenumber. By comparing this template with the CMB data over many different combinations of wavenumber it is possible to attain reasonable signal-to-noise. This comparison carries a considerable computational burden, so constraints from the data are typically reported as amplitudes for just a handful of well-known templates, such as the `local', `equilateral' and `orthogonal' shapes. These amplitudes are often written $\estfNLlocal$, $\estfNLequi$, $\estfNLortho$, and so on.% \footnote{Here and throughout the remainder of the paper we distinguish quantities estimated from data by a hat.} A specific inflationary model will be characterized by a number $\Nparams$ of adjustable parameters $\lambda_i$, $1 \leq i \leq \Nparams$. These may include Lagrangian parameters which are analogues of masses and couplings, but in multiple-field models may also include a specification of the initial conditions in field-space. To apply constraints from $\estfNLlocal$, $\estfNLequi$, $\estfNLortho$, \ldots, to such a model its three-point function must be computed and projected on to each of these templates. This generates predictions for each of the amplitudes $\fNLlocal(\lambda_i)$, $\fNLequi(\lambda_i)$, $\fNLortho(\lambda_i)$, \ldots. The results obtained by a microwave background survey can then be converted into constraints on the underlying parameters $\lambda_i$. This approach is perfectly reasonable, but there are reasons to expect that it may not be optimal. First, if the set of templates does not cover the entire range of three-point correlations which can be produced by adjusting the parameters $\lambda_i$ then we are not making efficient use of the data: we should measure the amplitude of more templates in order to obtain better constraints. But, as many authors have pointed out, it is not clear \emph{a priori} how large a range of templates is required, or how they should be chosen. Second, if our templates are chosen injudiciously then there will come a point of diminishing returns at which no new information is gained because the shapes we are fitting are strongly correlated with shapes which have been tried before. This is a reflection of a more general problem: the error bars reported for any set of amplitudes will typically be correlated, with the correlation described by some covariance matrix. Without knowledge of these covariances we risk underestimating the uncertainties associated with our reconstruction of the parameters $\lambda_i$. In this paper we take a different approach. We investigate the construction of maximum-likelihood estimators for the Lagrangian parameters $\lambda_i$ directly from the data. (Because noise maps for the Planck2013 data release are not yet available, we use the WMAP 9-year dataset.) To decide which templates to use, we catalogue the different types of correlation which can be produced in a well-specified class of models: those whose fluctuations are described the the effective field theory of inflation~\cite{Cheung:2007st}. We construct the Fisher matrix associated with these correlations and use it to determine the principal directions whose amplitudes can be measured efficiently. We account for the covariance between measurements of these amplitudes and use them to place constraints on the underlying Lagrangian parameters. \para{Summary}% In~{\S\ref{sec:EFT}} we briefly review the effective field theory approach to single-field inflation and catalogue the operators arising from a general single-field action. In~{\S\ref{sec:bispec}} we discuss the calculation of bispectra corresponding to these operators, and point out a number of subtleties which must be borne in mind when interpreting our results. In~{\S\ref{sec:estimating}} we assemble the formalism which is used to extract constraints from the CMB map: in~{\S\ref{sec:how-many-shapes}} we construct the Fisher matrix and use it to determine the principal directions which can be constrained efficiently, and in~{\S\ref{sec:results}} we report our measurements of their amplitudes from the 9-year WMAP dataset. {\S\ref{sec:model-constraints}} translates these general constraints into the language of specific models, and~{\S\ref{sec:model-comparison}} uses the framework of Bayesian model comparison to gain some qualitative information regarding the type of model favoured by the data. We conclude in~{\S\ref{sec:conclusions}}. A short appendix tabulates the three-point functions used in the main text. \para{Notation}% We use units in which $c = \hbar = 1$, and define the reduced Planck mass $\Mp$ to be $\Mp^{-2} = 8 \pi G$. Our index and summation conventions are explained in the main text.
\label{sec:conclusions} The availability of high-quality maps of the CMB temperature anisotropy from the WMAP and Planck missions means that it has become feasible to search for primordial three-point correlations. Such correlations are typically predicted by any scenario in which the fluctuations have an inflationary origin, due to microphysical three-body interactions among the light, active degrees of freedom of the inflationary epoch. If detected, their precise form could provide decisive evidence in favour of the inflationary hypothesis. Unfortunately, due to issues of computational complexity, it is not yet possible to perform a blind search for these primordial three-point correlations. Instead, we must search for signals which we have some prior reason to believe may be present in the data. Therefore the amount of information we manage to extract depends on which signals we choose to look for. In this paper we have made a systematic search of the 9-year WMAP data for correlations which could be produced in a very general model of single-field inflation, under the assumption that the background evolution is smooth, yielding corresponding smooth and nearly scale-invariant correlation functions. This excludes models which contain sharp features or oscillations~\cite{Starobinsky:1992ts, Adams:1997de,Adams:2001vc,Hailu:2006uj, Bean:2008na,Achucarro:2010da,Joy:2007na, Hotchkiss:2009pj,Nakashima:2010sa,Adshead:2011bw}. It also excludes models in which significant three-point correlations are generated by differences of evolution between regions of the universe separated by super-Hubble distances. Correlations generated by this mechanism are generally most significant in the `squeezed' or soft limit, where the correlation is between fluctuations on very disparate scales. Such correlations have been disfavoured by analysis of the Planck2013 data release~\cite{Ade:2013ydc}. By comparison, the 9-year WMAP data achieve a smaller signal-to-noise for such configurations. The difference between the 9-year WMAP and Planck2013 datasets is less pronounced for the momentum configurations which we probe, with for example $1\sigma$ error bars on $\fNL^{\rm{equil}}$ improving from $117$ to $75$. The essential steps of our analysis were assembled in~{\S\S\ref{sec:EFT}--\ref{sec:estimating}}. We begin with an effective field theory which parametrizes the unknown details of three-body interactions between inflaton fluctuations, but preserves nonlinearly realized Lorentz invariance. The effective theory is agnostic regarding the physical mechanism which underlies inflation. We compute the bispectrum generated by each operator in the effective theory, and break these into principal components using a Fisher-matrix approach. The amplitude of each principal component is recovered from the data, after which the results can be translated into constraints on the mass scales which appeared in the original effective theory. We find that no significant deviation from Gaussianity has been detected in any region of the inflationary parameter space. This conclusion is consistent with previous analyses of the 9-year WMAP and Planck2013 datasets. Our principal components are similar to those obtained by Byun \& Bean, who forecast the constraints which could be obtained from a Planck-like survey~\cite{Byun:2013jba}. We find that the best-constrained principal direction exhibits similarities to (in order) the flattened, orthogonal and `Galileon' templates. A fourth principal direction is more complex, but at best weakly constrained. The large space of models which fit into the class of single-field scenarios invites attempts to identify best-fitting regions. To approach this problem we use the framework of Bayesian model comparison. The results are at best weakly significant, but tend to disfavour models with more parameters when compared to simpler cases with zero or one parameter. This is not surprising given that the amplitude of each principal direction is consistent with zero. However, it should be borne in mind that our analysis is restricted to smooth and nearly scale-invariant bispectra. It is possible that a significant signal of a different type is hidden in the data. In some cases, $n$-point functions of this type can be described within the framework of effective field theory~\cite{Bartolo:2013exa}. The analysis developed in~{\S\S\ref{sec:EFT}--\ref{sec:estimating}} could be applied immediately to such scenarios given a suitable choice of basis functions $\Rmode_n$.
14
3
1403.3403
We use WMAP 9-year bispectrum data to constrain the free parameters of an `effective field theory' describing fluctuations in single-field inflation. The Lagrangian of the theory contains a finite number of operators associated with unknown mass scales. Each operator produces a fixed bispectrum shape, which we decompose into partial waves in order to construct a likelihood function. Based on this likelihood we are able to constrain four linearly independent combinations of the mass scales. As an example of our framework we specialize our results to the case of `Dirac-Born-Infeld' and `ghost' inflation and obtain the posterior probability for each model, which in Bayesian schemes is a useful tool for model comparison. Our results suggest that DBI-like models with two or more free parameters are disfavoured by the data by comparison with single-parameter models in the same class.
false
[ "model comparison", "unknown mass scales", "Bayesian schemes", "single-parameter models", "comparison", "single-field inflation", "Bayesian", "DBI-like models", "partial waves", "fluctuations", "the free parameters", "the mass scales", "each model", "the same class", "order", "a useful tool", "operators", "the posterior probability", "an `effective field theory", "a likelihood function" ]
11.486852
-0.506631
89
885812
[ "Petrie, G. J. D." ]
2014SoPh..289.3663P
[ "Estimating Flare-Related Photospheric Lorentz Force Vector Changes Within Active Regions" ]
11
[ "National Solar Observatory, 85719, Tucson, AZ, USA" ]
[ "2015ApJ...806..173B", "2015RAA....15..145W", "2016NatCo...713798B", "2016RAA....16...95Y", "2016RAA....16..129K", "2016SoPh..291..791P", "2018ApJ...869...21L", "2019ApJS..240...11P", "2019ApJS..242...30P", "2019LRSP...16....3T", "2020MNRAS.497..976K" ]
[ "astronomy" ]
7
[ "Active regions", "magnetic fields", "Flares", "dynamics", "relation to magnetic field", "Magnetic fields", "photosphere", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1964ARA&A...2..363S", "1969PhDT.........3H", "1973asqu.book.....A", "1974SoPh...39..393M", "1981ApJ...250..791Z", "1985svmf.nasa...49L", "1992PASJ...44L..89M", "1992SoPh..142...11W", "1994ApJ...424..436W", "1995ApJ...439..474M", "1999ApJ...515..842A", "1999SoPh..186..123G", "1999SoPh..190..459K", "2000ApJ...531L..75H", "2001ApJ...550L.105K", "2001SoPh..203...71G", "2004ApJ...613L.177L", "2005ApJ...635..647S", "2005JGRA..110.8104Z", "2006ApJ...649..490W", "2008ASPC..383..221H", "2008SoPh..247..411T", "2010ApJ...716L.195W", "2010ApJ...724.1218P", "2010JASTP..72..219H", "2011SoPh..271...27F", "2012ApJ...745L...4L", "2012ApJ...745L..17W", "2012ApJ...748...77S", "2012ApJ...749...85G", "2012ApJ...759...50P", "2012SoPh..275....3P", "2012SoPh..275..229S", "2012SoPh..277...59F", "2012SoPh..280..335A", "2013SoPh..287..415P", "2013SoPh..288..617R" ]
[ "10.1007/s11207-014-0545-3", "10.48550/arXiv.1403.6156" ]
1403
1403.6156_arXiv.txt
\label{sect:intro} In recent years it has repeatedly been shown that the solar photosphere responds to the abrupt coronal reconfiguration of most major flares, often with abrupt, coherent, permanent, and widely distributed patterns of photospheric magnetic field change. Flare-related photospheric magnetic field changes were sought for many years ({\it e.g.}, Severny,~1964; Zirin and Tanaka,~1981), before abrupt, permanent photospheric field changes were successfully observationally linked to flares; see the discussion in Sudol and Harvey~(2005). Since the launch of NASA's {\it Solar Dynamics Observatory} (SDO) satellite (Pesnell {\it et al.},~2012), the {\it Helioseismic and Magnetic Imager} (HMI) instrument (Schou {\it et al.},~2011) has provided full-disk vector magnetograms with $0.\!\!^{\prime\prime}5$ pixels every 12 min, enabling the study of flare-related photospheric vector magnetic field changes in more detail than was possible previously. The history of flare-related photospheric magnetic field changes has not been straightforward. Following the groundbreaking work of Wang {\it et al.}~(1992, 1994), reporting abrupt, permanent field changes in flaring active regions, a number of later studies produced inconclusive results; see the discussion in Wang~(2006). Then, the introduction of high-cadence longitudinal (line-of-sight) photospheric magnetic field measurements, from NASA's {\it Michelson Doppler Imager} (MDI) on the {\it Solar and Heliospheric Observatory} (SoHO) satellite and the National Solar Observatory’s {\it Global Oscillations Network Group} (GONG) ground-based network, resulted in mounting evidence of flare-related magnetic changes in the photosphere. Based on one-min MDI longitudinal data, Kosovichev and Zharkova~(1999, 2001) reported sudden and permanent decreases in magnetic flux near X-class flares near magnetic neutral lines, and Zharkova {\it et al.}~(2005) found a permanent and significant increase in magnetic flux near the neutral line of a region during another X-class flare observed near the limb. Using one-min GONG longitudinal magnetograms, Sudol and Harvey~(2005) characterized the spatial distribution, strength, and rate of change of permanent field changes associated with 15 X-class flares. By carefully co-registering the images they succeeded in tracing the field changes pixel-by-pixel and were able to show the spatial structure of the changes. They concluded that the majority of field changes occurred close to or within sunspots. Wang~(2006) found an unshearing movement parallel to the neutral lines in flare-related longitudinal magnetic field changes in the MDI data for all five $\delta$-spot flares that he studied, implying an overall release of shear, but that the two polarities converged toward the neutral line during some events and diverged during others. Petrie and Sudol~(2010) analyzed 1-min GONG longitudinal magnetograms covering 77 flares of GOES class at least M5 and, exploring the relationship between increasing/decreasing longitudinal fields and azimuth and tilt angles at various positions on the disk, found that the overall distributions of longitudinal increases and decreases at different parts of the disk was found to be consistent with the coronal implosion interpretation (Hudson,~2000; Hudson, Fisher, and Welsch,~2008; Fisher {\it et al.},~2012) where, after a coronal magnetic eruption, the remaining coronal field loop structure contracts downward resulting in a more horizontal photospheric field. Wang and Liu~(2010) studied 11 X-class flares using vector magnetograms from the vector magnetograph of Big Bear Solar Observatory and other vector magnetographs, and found in each case an increase of transverse field at the main polarity-inversion line. The launch of SDO and the release of 12-min HMI photospheric vector magnetograms prompted many authors to study flare-related photospheric vector field changes. Wang {\it et al.}~(2012), Gosain~(2012), Sun {\it et al.}~(2012) and Petrie~(2012, 2013) analyzed the 12-min HMI vector data for the 15 February 2011 X2.2 flare using different methods, and again found an increase of transverse field at the polarity-inversion line, as did Liu {\it et al.}~(2012) for the 13 February 2011 M6.6 flare. These observations support Hudson, Fisher, and Welsch’s~(2008) loop-collapse scenario. Hudson, Fisher, and Welsch~(2008) and Fisher {\it et al.}~(2012) developed a method for estimating the Lorentz force change acting on the photosphere from the atmosphere above using the changes in the photospheric vector field observed during (or shortly before and after) the flare. The Lorentz force is estimated by integrating components of the Maxwell stress tensor across the boundary of a large domain containing the active region (see Figure~1 of Fisher {\it et al.},~2012). The bottom boundary of the domain is identified with the photospheric layer from which the observations derive, and the field values on this boundary are represented by the photospheric magnetic field observations. Because the field values at the lateral and top boundaries of the domain are not measured, Fisher {\it et al.}~(2012) advocated applying their method only to domains where these boundaries are far enough away from the region that the fields crossing them are too weak to contribute significantly to the integral, allowing one to ignore them. Strict adherence to this approach eliminates the possibility of studying the diverse responses of distinct components of an active region, such as sunspots and neutral-line arcades, whose behavior may provide important clues regarding the physics of the flare. Because of this, some authors have proceeded to apply Fisher {\it et al.}'s~(2012) equations to subdomains within active regions ({\it e.g.}, Wang {\it et al.},~2012; Petrie,~2012, 2013), sometimes providing brief arguments in favor of doing so (Alvarado-G\'omez {\it et al.},~2012; Petrie,~2013). In particular, it has been argued that the force-free nature of the coronal magnetic field allows one to neglect most of the contribution from the lateral and top boundaries of such a domain. In this paper we will revisit this problem using a model for the stratified solar atmospheric field. When the 12-min HMI vector data became available it became possible to apply Fisher {\it et al.}'s~(2012) surface integrals for the photospheric Lorentz force, and it is noteworthy that these calculations have produced strikingly consistent results. Wang {\it et al.}~(2012) found that the vertical Lorentz force changes associated with 18 major flares were generally directed downward near neutral lines and were correlated with peak soft X-ray flux. Petrie~(2012) studied abrupt changes in both the photospheric magnetic and Lorentz force vector associated with six major neutral-line flares using HMI vector data. During all six flares the neutral-line field vectors became stronger and more horizontal, in each case almost entirely due to strengthening of the horizontal field components parallel to the neutral line. In all cases the neutral-line pre-flare fields were less tilted than the reference potential fields, and collapsed abruptly and permanently closer to potential-field tilt angles, implying that the relaxation of magnetic stress played a leading role in creating the magnetic changes. Indeed, the vertical Lorentz force had a large, abrupt, permanent downward change at the main neutral line during each of the flares, consistent with loop collapse. The horizontal Lorentz force changes acted mostly parallel to the neutral line in opposite directions on each side, a signature of the fields contracting during this loop collapse, pulling the two sides of the neutral line toward each other. The greater effect of the flares on field tilt than on shear may be explained by photospheric line-tying, since shear cannot be removed during loop collapse without the foot-points moving across the photosphere. Petrie~(2013) found that, in the case of the 15 February 2011 X2.2 flare, the oppositely-directed horizontal Lorentz force changes acting on each side of neutral line were accompanied by abrupt torsional un-twisting forces in the two sunspots at each end of the neutral line. While these studies exploited Fisher {\it et al.}'s (2012) equations to offer more physical insight into the flare-related magnetic changes than the magnetic vector information alone, it remained true that Fisher {\it et al.}~(2012) did not advocate applying their method to fields within active regions and that a detailed defense was lacking. The purpose of this paper is to investigate whether and under what circumstances Lorentz force estimates within active regions can be reliably produced using this method. In this paper we revisit Fisher {\it et al.}'s calculation in the context of Gary's~(2001) well-known model for the plasma $\beta$ (the ratio between the plasma and magnetic pressures) of an active region. This allows us to study the influence of solar atmospheric stratification on Lorentz-force distribution between the photosphere and the corona. The paper is organized as follows. In Section~\ref{sect:fisher} we review Fisher {\it et al.}'s~(2012) method. In Section~\ref{sect:subdomain} we derive expressions representing the contributions to the Lorentz force integral from the different boundary surfaces of a subdomain within an active region, and arrive at inequalities describing conditions where photospheric Lorentz forces within active regions can be derived from photospheric measurements alone. In Section~\ref{sect:gary} we introduce Gary's~(2001) plasma $\beta$ model and in Section~\ref{sect:lfintegral} we estimate force contributions from different atmospheric layers using this model. In Section~\ref{sect:conclusion} we arrive at conclusions regarding the application of Fisher {\it et al.}'s method to subdomains within active regions.
\label{sect:conclusion} \begin{figure} \begin{center} \resizebox{\textwidth}{!}{\includegraphics*{f4.eps}} \end{center} \caption{Vector magnetic field before the 15 February 2011 X2.2 flare. The vertical field component, $B_r$, is indicated by the color scale and the horizontal component by the arrows, with saturation values ±1000 G. Red/blue coloring represents positive/negative vertical field. The solid, dashed, dotted and dot-dashed contours indicate photospheric field strengths corresponding to the four models featuring in the previous figures and described in Section~\ref{sect:gary}. At the approximate height of the HMI observations, $z=100$~km, these models have field strength 1500~G, 629~G, 340~G, and 108~G, respectively.} \label{fig:feb15flare} \end{figure} We conclude that for strong magnetic fields the contribution to the Lorentz force estimate from fields at height above around 1~Mm can be neglected. For photospheric subdomains filled with large, coherent structure and strong magnetic field, the contributions to the surface integrals derived by Fisher {\it et al.}~(2012) are negligible compared to the contribution from the photospheric boundary. This implies that Fisher {\it et al.}'s~(2012) method for estimating Lorentz force changes on the photosphere during solar flares using Equations~(\ref{eq:fr}) and (\ref{eq:fh}) can be applied to subdomains within a flaring active region, so long as the conditions on strength of field and coherence of structure are met within this photospheric subdomain. For HMI data, and for data originating from similar heights in the solar atmosphere, the horizontal length scale of the structure needs to be much larger than 300~km and the field needs to be stronger than about 630~G at the height of the observations ($\ge 1$~kG at the $\tau =1$ opacity layer at 5000~\AA). This conclusion therefore supports many of the published surface integral calculations based on Fisher {\it et al.}'s method, such as those by Petrie~(2012, 2013), who used regular rectangular integration domains covering strong neutral line fields and circular domains covering sunspots, and Wang and Liu~(2010) and Wang {\it et al.}~(2012), who used irregular domains covering strong sunspot and neutral line fields. Figure~\ref{fig:feb15flare}\footnote{This figure is updated from Figure~1 of Petrie~(2012). In Petrie~(2012) the longitude-latitude coordinate system has origin at the center of the remapped HMI magnetogram whereas here the origin corresponds to disk-center. See also http://hmi.stanford.edu/hminuggets/?p=539} shows the much-studied HMI vector magnetic field associated with the 15 February 2011 X2.2 flare. The plot shows that most of the important fields at the heart of the active region NOAA 11158 are strong enough to meet the above conditions, and the major features of the region have horizontal spatial scale much larger than 300~km. Note that Fisher {\it et al.}'s~(2012) equations record only the Lorentz force difference between two vector magnetogram measurements covering the same locations at different times. During a highly dynamic event such as a flare, Lorentz forces apply to different layers of the solar atmosphere, but only those layers whose fields have significant and steady Lorentz force values before and after the flare will yield measurements that can be analyzed using this method. In a nearly force-free layer of the atmosphere the magnetic vectors may respond strongly to a flare. However, if the magnetic vector field in this layer is not accompanied by a significant and lasting Lorentz force change associated with the flare then Equations~(\ref{eq:deltafr}) and (\ref{eq:deltafh}) will not give a result that can be understood as a flare-related Lorentz force vector change. This is true even if the flare produced a large and permanent change in the magnetic vector field. In some nearly force-free layers of the solar atmosphere, dynamical equilibration processes remove any significant Lorentz forces from field configurations, erasing all traces of the Lorentz force changes produced by flares. It is therefore essential to the success of this method that the magnetogram observations derive from a layer where such force-removing equilibration processes do not generally occur. \begin{acks} I thank the referee for helpful comments. This project was prompted by discussions at Leverhulme Flare Seismology Workshop held at the Mullard Space Science Laboratory (MSSL), 9-12 September 2013. The author thanks the Leverhulme trust for hospitality during the author's visit to MSSL, and thanks his fellow participants at the workshop for interesting discussions during and after the meeting. \end{acks}
14
3
1403.6156
It is shown that expressions for the global Lorentz force associated with a flaring active region derived by Fisher et al. (Solar Phys.277, 59, 2012) can be used to estimate the Lorentz-force changes for strong fields in large structures over photospheric subdomains within active regions. Gary's (Solar Phys.203, 71, 2001) model for the stratified solar atmosphere is used to demonstrate that in large-scale structures with typical horizontal magnetic length scale ≫ 300 km and with strong magnetic fields (≥ 1 kG at the τ=1 opacity layer at 5000 Å), the Lorentz force acting on the photosphere may be approximated by a surface integral based on photospheric boundary data alone. These conditions cover many of the sunspot fields and major neutral lines that have been studied using Fisher et al.'s (2012) expressions over the past few years. The method gives a reasonable estimate of flare-related Lorentz-force changes based on photospheric magnetogram observations provided that the Lorentz-force changes associated with the flare have a lasting effect on the observed fields, and they are not immediately erased by post-flare equilibration processes.
false
[ "Fisher et al", "strong magnetic fields", "strong fields", "photospheric boundary data", "post-flare equilibration processes", "typical horizontal magnetic length scale", "active regions", "Lorentz", "photospheric magnetogram observations", "photospheric subdomains", "large structures", "Fisher et al.s", "the global Lorentz force", "major neutral lines", "the past few years", "the Lorentz force", "the Lorentz-force changes", "large-scale structures", "a flaring active region", "the observed fields" ]
12.602064
15.965299
2
812234
[ "Mangano, Gianpiero", "Mirizzi, Alessandro", "Saviano, Ninetta" ]
2014PhRvD..89g3017M
[ "Damping the neutrino flavor pendulum by breaking homogeneity" ]
45
[ "Istituto Nazionale di Fisica Nucleare - Sezione di Napoli, Complesso Universitario di Monte S. Angelo, I-80126 Napoli, Italy", "II Institut für Theoretische Physik, Universität Hamburg, Luruper Chaussee 149, 22761 Hamburg, Germany", "Institute for Particle Physics Phenomenology, Department of Physics, Durham University, Durham DH1 3LE, United Kingdom" ]
[ "2014PhRvD..90d5032T", "2014arXiv1405.7275A", "2015AIPC.1666g0001L", "2015IJMPE..2441008D", "2015NPPP..265...45M", "2015NPPP..265..333G", "2015PhLB..747..139D", "2015PhLB..751...43A", "2015PhRvD..92b1702M", "2015PhRvD..92f5019A", "2015PhRvD..92j5002C", "2015PhRvD..92j5020M", "2015PhRvD..92l5030D", "2016JCAP...01..028C", "2016JCAP...03..042C", "2016JCAP...04..043C", "2016JPhCS.718b2007D", "2016NCimR..39....1M", "2016NuPhB.908..366C", "2016PhLB..752...89W", "2017JCAP...02..019D", "2017PhRvD..96d3016C", "2017PhRvL.118b1101I", "2018JCAP...04..057H", "2018PhRvD..97b3017D", "2018PhRvD..98j3020Z", "2019JCAP...08..003D", "2019JCAP...10..027H", "2019PhRvD..99f3005Y", "2019PhRvD..99l3004B", "2019PhRvD.100b3016M", "2020PhRvD.101b3007M", "2020PhRvD.101d3009J", "2020PhRvD.101d3016A", "2020PhRvD.102f3018B", "2020arXiv200500459B", "2020arXiv201203948H", "2021JCAP...07..023B", "2021PhRvD.103d5014A", "2021PhRvD.103f3008Z", "2021PhRvL.126f1302B", "2022PhRvD.105d3005S", "2022PhRvD.105l3036M", "2022Univ....8...94C", "2023PhRvD.107j3022Z" ]
[ "astronomy", "physics" ]
9
[ "14.60.Pq", "97.60.Bw", "Neutrino mass and mixing", "Supernovae", "High Energy Physics - Phenomenology", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
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[ "10.1103/PhysRevD.89.073017", "10.48550/arXiv.1403.1892" ]
1403
1403.1892_arXiv.txt
Neutrino--neutrino interactions in dense neutrino media are known to produce surprising flavor oscillation effects, in the form of self-induced conversions, when the typical neutrino self-interaction potential $\mu = {\sqrt{2} G_F}n_\nu$ is comparable with or greater than the vacuum oscillation frequency $\omega=\Delta m^2/2E$ (see e.g.~\cite{Duan:2010bg} for a recent review). This situation can be encountered in the early universe or in core-collapse supernovae (SN), where neutrino themselves form a background medium for their propagation. Differently from the usual Mykheyeev-Smirnov-Wolfenstein (MSW) effect~\cite{Matt}, associated with the matter potential $\lambda=\sqrt{2}G_F n_e$, the self-induced effects do not change the flavor content of the neutrino ensemble. Yet, the flavor is exchanged between different momentum modes, leading to peculiar spectral features known as {\it spectral swap} and {\it split}~\cite{Raffelt:2007cb}. The growth of these effects is associated with instabilities in the flavor space, which are amplified by the neutrino-neutrino interactions~\cite{Sawyer:2005jk,Banerjee:2011fj}. An example is represented by the bimodal instability~\cite{Samuel:1995ri} of an isotropic and homogeneous dense gas of neutrinos and antineutrinos in equal amounts. They convert from one flavor to another in pair production processes $\nu_e \bar\nu_e \leftrightarrow \nu_x \bar\nu_x$, behaving as a flavor pendulum even if the mixing angle is very small~\cite{Hannestad:2006nj,Duan:2007fw}. In this case, the vacuum mixing angle acts as a seed triggering the flavor instability. In non-isotropic neutrino gases, like the case of neutrinos streaming-off a SN core, the features of the self-induced effects are more involved, since the current-current nature of the low-energy weak interactions introduces an angle dependent term $(1-{\bf v}_{\bf p} \cdot {\bf v}_{\bf q})$, where ${\bf v}_{\bf p}$ and ${\bf v}_{\bf q}$ are neutrino velocities~\cite{Qian:1994wh,Duan:2006an}. It has been shown that this term can lead to a {\it multi-angle} instability, which hinder the maintenance of the coherent oscillation behavior for different neutrino modes~\cite{Duan:2006an,EstebanPretel:2007ec,Sawyer:2008zs}. In particular, in a symmetric gas of equal neutrino and antineutrino densities even a very small anisotropy is sufficient to trigger a run-away towards flavor equipartition~\cite{decoh}. An additional instability has been recently discovered in the SN context. Removing the assumption of axial symmetry in the $\nu$ propagation, a multi-azimuthal-angle instability emerges, even assuming a perfect spherically symmetric $\nu$ emission~\cite{Raffelt:2013rqa,Raffelt:2013isa,Duan:2013kba,Mirizzi:2013rla,Mirizzi:2013wda,Chakraborty:2014nma}. Symmetries in the neutrino self-induced evolution are often assumed in order to reduce the complexity of the problem. Nevertheless, these recent findings {\it question} the validity of these assumptions, since they suggest that (unavoidable) small deviations from initial symmetries could be dramatically amplified by the interacting neutrinos during the evolution. In absence of collisions, the dynamics of the $\nu$ space-dependent occupation numbers or Wigner function $\varrho_{{\bf p}, {\bf x}}(t)$ with momentum ${\bf p}$ at position ${\bf x}$ is ruled by the kinetic equations~\cite{Sigl:1992fn,Strack:2005ux} \begin{eqnarray} && \partial_t \varrho_{{\bf p}, {\bf x}} + {\bf v}_{\bf p} \cdot \nabla_{\bf x} \varrho_{{\bf p}, {\bf x}} + {\dot{\bf p}} \cdot \nabla_{\bf p} \varrho_{{\bf p}, {\bf x}} \nonumber \\ && = - i [\Omega_{{\bf p}, {\bf x}}, \varrho_{{\bf p}, {\bf x}}] \,\ , \label{eq:eom} \end{eqnarray} with the Liouville operator in the left-hand side. In particular, the first term accounts for an explicit time dependence, while the second is the drift term proportional to the neutrino velocity ${\bf v_p}$, due to particle free streaming. Finally, the third term is proportional to the force acting on neutrinos. On the right-hand-side the matrix $\Omega_{{\bf p}, {\bf x}}$ is the full Hamiltonian containing the vacuum, matter and self-interaction terms. We remind the reader that the quantum-mechanical uncertainty between location and momentum implies that this formalism can be applied only for cases where spatial variations of the ensemble are weak on the microscopic length-scale defined by the typical particle wave-length. In general, Eq.~(\ref{eq:eom}) describes a seven-dimensional problem that has never been solved in its complete form. For neutrino flavor conversions in the early universe one typically assumes initial \emph{space} homogeneity, that allows one to reduce the dependence on space-time variables to time only. Conversely, for neutrinos in a SN environment, a spatial evolution under the assumption of a \emph{stationary} neutrino emission is often considered. For a spherically symmetric neutrino emission with negligible variations in the transverse direction, the description further simplifies, the problem being reduced to a purely radial dynamics. However, small space inhomogeneities over the standard rotation and translation invariant background are expected in the early universe, with an initial spectrum in Fourier space very close to the scale invariant Harrison-Zeldovich one, the typical heritage of an inflationary expansion initial stage. On the other hand, in the SN environment one should account for deviations with respect to a stationary configuration, which are related to hydrodynamical instabilities. Both these deviations with respect to the assumed homogeneity conditions can act as seeds for instabilities. In order to investigate this issue, rather than studying the behavior of the complex early universe or SN systems, we consider here a much more simple toy model, which however already illustrates the main point of this paper. Namely, unless spatial symmetry (or stationarity) is imposed by hand, the self-interacting neutrino dynamics is unstable with respect to even tiny ripples over a spatially constant (or time independent) background. In particular, we consider a neutrino ensemble in time and one spatial dimension, initially prepared with equal densities of $\nu_e$ and ${\bar\nu}_e$ evolving in time. In the position invariant case, this system behaves as a flavor pendulum in inverted mass hierarchy. We then introduce a small space-dependent fluctuation in the matter potential and look for its effect on the evolution of the neutrino density matrix. We stress that in our case the matter inhomogeneities just act as a seed to trigger a possible instability and are chosen to have amplitudes much smaller than those often considered in literature in relation to the MSW effect (see, e.g.,~\cite{Sawyer:1990tw,Kneller:2012id}). The paper is organized as follows. In Section II we describe the equations of motion for the neutrino ensemble evolving in time in presence of inhomogeneities. By Fourier transforming the equations in the spatial coordinate, the problem is then reduced to ordinary differential equations in the time variable for the different Fourier modes, which are coupled each other. In Section III we numerically solve these equations for a constant background neutrino potential $\mu > \omega$ and we show how the system exibits a run--away from the flavor pendulum behavior, even for very small matter perturbations. The decoherence is indeed, associated with the growth of the different Fourier modes that destabilize the ordered pendulum solution. We also consider the effect of a time depending neutrino self-potential on this instability. If $\mu$ is a decreasing function, as we expect to be the case in both the early universe (with respect to time) or SN scenarios (with respect to distance), the growth of higher wave number Fourier modes might be inhibited, for a sufficiently short decay time scale. Finally, in Section IV we summarize our results, we comment about the possible effects in realistic neutrino gases and we conclude.
The role of symmetries in reducing the complexity of the dynamics of self--interacting neutrino systems has been widely exploited. The general seven--dimensional differential problem can be reduced to more treatable models, and numerically solved with less demanding computation powers. However, simplifying the scenarios to the only time or radial evolution of neutrino density matrix, neglecting space inhomogeneities or non stationary features, provides useful results only if the dynamics is {\it stable} against perturbations. If this is not the case, the behavior of neutrino medium can be very different from what is found when a particular symmetry is imposed by hand. In this context we have investigated the emergence of a new kind of instability in the flavor evolution of a dense neutrino gas, when space homogeneity is slightly perturbed. In order to illustrate this effect we have considered the time evolution of a simple system based on an isotropic neutrino ensemble, initially composed by only $\nu_e$ and $\bar\nu_e$ with equal densities. We have introduced a small amplitude position-dependent perturbation. In order to follow the simultaneous temporal and spatial flavor evolution we have Fourier transformed the equations of motion, obtaining a tower of equations for the different modes associated with the space variable. These modes are coupled because of the effect of the inhomogeneities on the neutrino-neutrino interaction term. We found that in inverted mass hierarchy, where an homogeneous neutrino gas would have evolved according to the flavor pendulum solution~\cite{Hannestad:2006nj}, the presence of the inhomogeneous term destroys this coherent behavior and leads to flavor decoherence. This is due to the growing of modes with non zero wavenumber, excited by their coupling to the neutrino homogeneous zero mode. The system is instead stable in the normal mass hierarchy case. The instability discussed here complements the findings of~\cite{decoh}, where for the same $\nu-\bar\nu$ symmetric case, breaking the isotropy of the neutrino propagation leads to a quick decoherence in both mass hierarchies. In that case the equations of motion were expanded in multipoles through the Legendre functions and the decoherence was associated to the excitement of the higher multipoles. Breaking the homogeneity we are observing a similar phenomenology. Despite the simplicity of our model -- one space dimension and a monochromatic matter potential disturbance -- we think that features similar to those described in this paper would be also present in more realistic cases. There are indeed, two main frameworks where inhomogeneities (in time or space) could play a role in self-induced flavor conversions: neutrinos evolution in the early universe or in their streaming off core--collapse supernovae. In the first case we know that lepton and neutrino number densities keep the imprint of small space perturbations over the homogeneous background configuration. These are likely produced during the inflationary phase with an almost scale invariant spectrum and initial amplitudes of order $10^{-5}$. In the second example, deviations from a stationary evolution are expected to be triggered by hydrodynamical instabilities. In both these ensembles the neutrino density decreases with respect to the evolution variable. In order to mimic this effect, we have considered a declining matter potential $\mu$. We found that this could inhibit the growth of the different Fourier modes, since the evolution is not enough adiabatic to let them to develop. The decoherence effect could be thus, suppressed in realistic scenarios. However, it is not guaranteed that the simultaneous breaking of isotropy and homogeneity may not lead to novel phenomena. At this regard the impact of multi-angle effects requires a detailed investigation that we leave for a future work. Finally, we would like to stress that the formalism we have developed here to treat the simultaneous time and space flavor evolution could be applied to study other deviations from a stationary SN neutrino flavor evolution, as those which could be induced by the small backward flux caused by residual neutrino scattering that causes significant refraction~\cite{Cherry:2012zw,Sarikas:2012vb}. In any case it is intriguing that even the simplest neutrino flavor pendulum is still a source of new instabilities that were not appreciated before. \\
14
3
1403.1892
The most general case of self-induced neutrino flavor evolution is described by a set of kinetic equations for a dense neutrino gas evolving in both space and time. Solutions of these equations have been typically worked out assuming that either the time (in the core-collapse supernova environment) or space (in the early Universe) homogeneity in the initial conditions is preserved through the evolution. In these cases, one can gauge away the homogeneous variable and reduce the dimensionality of the problem. In this paper, we investigate whether small deviations from an initial postulated homogeneity can be amplified by the interacting neutrino gas, leading to a new flavor instability. To this end, we consider a simple two-flavor isotropic neutrino gas evolving in time, and initially composed by only ν<SUB>e</SUB> and ν ¯e with equal densities. In the homogeneous case, this system shows a bimodal instability in the inverted mass hierarchy scheme, leading to the well-studied flavor pendulum behavior. This would lead to periodic pair conversions ν<SUB>e</SUB>ν ¯e↔ν<SUB>x</SUB>ν ¯x. To break space homogeneity, we introduce small amplitude space-dependent perturbations in the matter potential. By Fourier transforming the equations of motion with respect to the space coordinate, we then numerically solve a set of coupled equations for the different Fourier modes. We find that even for arbitrarily tiny inhomogeneities, the system evolution runs away from the stable pendulum behavior: the different modes are excited and the space-averaged ensemble evolves towards flavor equilibrium. We finally comment on the role of a time decaying neutrino background density in weakening these results.
false
[ "space homogeneity", "flavor equilibrium", "pendulum behavior", "kinetic equations", "coupled equations", "equal densities", "neutrino background density", "self-induced neutrino flavor evolution", "time", "ν", "small amplitude space-dependent perturbations", "a new flavor instability", "the different Fourier modes", "the interacting neutrino gas", "small deviations", "the space coordinate", "the space-averaged ensemble evolves", "the stable pendulum behavior", "periodic pair conversions", "the different modes" ]
5.806692
-1.521252
45
746700
[ "Castro, T.", "Quartin, M." ]
2014MNRAS.443L...6C
[ "First measurement of σ8 using supernova magnitudes only." ]
20
[ "Instituto de Física, Universidade Federal do Rio de Janeiro, CEP 21941-972 Rio de Janeiro, RJ, Brazil", "Instituto de Física, Universidade Federal do Rio de Janeiro, CEP 21941-972 Rio de Janeiro, RJ, Brazil" ]
[ "2015ApJ...806...45Z", "2015JCAP...01..008Q", "2015JCAP...12..033H", "2015MNRAS.449.2845A", "2016ApJ...828..112H", "2016MNRAS.455..552B", "2016MNRAS.463.1666C", "2016arXiv161009452D", "2017JCAP...05..015H", "2017MNRAS.465.2862S", "2017MNRAS.467..259M", "2017hsn..book.2623D", "2018JCAP...07..024D", "2018PhRvL.121n1101Z", "2019PhRvD..99h3525Z", "2020MNRAS.496.4051M", "2020MNRAS.497.2974V", "2022LRR....25....6M", "2022MNRAS.512.2841Q", "2024arXiv240500147A" ]
[ "astronomy" ]
6
[ "gravitational lensing: weak", "cosmological parameters", "cosmology: observations", "distance scale", "large-scale structure of Universe", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
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[ "10.1093/mnrasl/slu071", "10.48550/arXiv.1403.0293" ]
1403
1403.0293_arXiv.txt
\label{sec:intro} Type Ia supernovae (SNeIa) are arguably the most important and reliable estimators of extragalactic distances. As it is well know, they provided the first solid evidence of the present cosmological acceleration~\citep{Riess:1998cb,Perlmutter:1998np}. Since this discovery a large effort has been devoted to testing and improving the calibration of the SNeIa and to correcting their light curves in order to understand and control systematics~\citep{Kessler:2009ys,Conley:2011ku,Betoule:2012an,Scolnic:2013aya}. As their light comes from high redshifts (up to $z\simeq2$) gravitational lensing from intervening matter is expected to play an important role. The correction induced by lensing will in fact become a major source of uncertainty when richer and deeper SNeIa catalogs are compiled in the next years. The Large Synoptic Survey Telescope (LSST) project plans for instance to collect over a million SNeIa in ten years~\citep{Abell:2009aa}, roughly a thousand-fold increase from number of SNeIa observed so far. A great effort is therefore being put forward to better understand this and avoid biases; see e.g.~\citep{Jonsson2008,Amendola:2010ub,Takahashi:2011qd,Clarkson:2011br,Bolejko:2012ue,BenDayan:2013gc,Zitrin:2013jza}. Gravitational lensing changes the intrinsic distribution function of the SNeIa magnitudes, increasing the scatter and introducing non-Gaussianity. In~\citep{Amendola:2013twa}, we have obtained the lensing variance, skewness and kurtosis of the SNeIa distribution via sGL, a fast simulation method developed in~\citep{Kainulainen:2009dw,Kainulainen:2010at,Kainulainen:2011zx}. When confronted to $N$-body simulations sGL was shown to be very accurate up to $z \simeq 1.5$, with the advantage of results being given as function of the relevant cosmological parameters. They also were in very good agreement with observational data~\citep{Jonsson:2009jp,Kronborg:2010uj,Jonsson:2010wx} and with other recent independent theoretical estimations~\citep{BenDayan:2013gc}. These fits can be employed to take into account the lensing extra scatter for any value of the cosmological parameters and also to model the lensing non-Gaussianity. This fact was explored in~\citep{Quartin:2013moa} where we proposed to use these accurate determinations of the lensing moments to measure cosmological parameters, following the ideas first discussed in~\citep{Bernardeau:1996un,Hamana:1999rk,Valageas:1999ir} and later further developed in~\citep{Dodelson:2005zt}. We showed that by using not just the variance of the lensing signal but the third and fourth order moments as well, a more precise and robust measurement was possible. In a $\Lambda$CDM scenario it was verified that the most sensitive cosmological parameters to supernova lensing were $\Omega_{m0}$ and $\sigma_8$. Now since the former is already tightly constrained by the measurement of the supernova magnitudes themselves (i.e., by the first moment of the distribution), the most important \emph{new} information gained was that pertaining to $\sigma_8$. In particular it was shown that $\sigma_8$ could be measured by the LSST survey to within 3--7\%, a value that is competitive with usual methods based on cosmic shear, cosmic microwave background (CMB) or cluster abundance, and completely independent of these. In particular, it does not rely on measuring galaxy shapes and is thus immune to the systematics associated to the cross-correlation of intrinsic galaxy ellipticities. Also, it does not require to extrapolate the amplitude $\sigma_8$ from recombination epoch to today, as with the CMB technique, nor to make assumptions on the threshold of formation of structures that is needed when employing galaxy clusters. It also complements the method % proposed in~\citep{Gordon:2007zw}, to wit correlating nearby supernova magnitudes with their positions to obtain their peculiar velocity correlations, which is also sensitive to $\sigma_8$. Here we extend on previous works on two fronts. First, we generalize the method to include intrinsic non-Gaussianities in the SNeIa distributions (that is, excluding all lensing effects). We do so by employing one nuisance parameter for each central moment of the distribution. We then argue that this is the most straightforward extension of the standard supernova analysis and that a more complicated parametrization should only be used if data itself demands it; the Bayes Factor is a nice and simple way to decide which parametrization to use. Second, we apply the above generalized procedure to two real supernova catalogs: the recently published combined SDSS-II and SNLS 3-year results~\citep{Betoule:2014frx}, dubbed the Joint Lightcurve Analysis (JLA) catalog and the older standard SNLS 3-year catalog (SNLS3)~\citep{Conley:2011ku}. We find that the method works as is, even though data is usually not treated for systematics that affect the higher moments. We thus obtain the first measurement of $\sigma_8$ from supernova magnitudes alone. This letter is organized as follows. In Section~\ref{sec:moments} we summarize our methodology. In Section~\ref{sec:bayes} we show how the Bayed Factor can be used to best model the SNeIa probability distribution function (PDF), and in Section~\ref{sec:snls} we apply our method to real data. Finally, we conclude in Section \ref{sec:conclusions}.
\label{sec:conclusions} In this letter we obtained the first constraints for $\sigma_8$ from SNeIa data alone. In other words, without need to cross-correlate SNeIa with matter distribution data, as done for instance in~\citep{Smith:2013bha}. In order to obtain such bounds we used two nuisance parameters to cope with intrinsic scatter and skewness in the data. In principle one can use also a third nuisance parameter for the kurtosis, but data showed no need of it. In fact, for the JLA catalog even $\mu_{3,{\rm int}}$ could be set to zero, but we chose to leave it and marginalize over to get more conservative results. Nevertheless, although the obtained bounds for $\sigma_8$ are very broad and systematics may be present, the consistency of the data with our mocks serves as an important validation of the method and opens up a new avenue in cosmology. In the future in order to best use this lensing information it is important to study whether experimental details or data reduction methods introduce systematics in the form of non-Gaussianities. Moreover, here we made use of the inferred SNeIa distances directly from JLA and SNLS3 catalogs. It would be interesting to check in detail whether including the $\sigma_8$ dependence due to lensing in the lightcurve fitter itself (i.e., simultaneously with the stretch and color corrections) significantly affects any of the results. It is clear that other similar tests can be employed with our methods. For instance, one can fix completely the cosmology at, say, the CMB values and just do a hypothesis test on the data as a consistency check with lensing predictions. Other interesting possibilities would be instead to use SNeIa lensing to test either the power spectrum directly~\citep{Ben-Dayan:2013eza} or the halo models~\citep{Fedeli:2013yfa}, but both require re-deriving our estimates for the central moments.
14
3
1403.0293
A method was recently proposed which allows the conversion of the weak-lensing effects in the Type Ia supernova (SNeIa) Hubble diagram from noise into signal. Such signal is sensitive to the growth of structure in the universe, and in particular can be used as a measurement of σ<SUB>8</SUB> independently from more traditional methods such as those based on the cosmic microwave background, cosmic shear or cluster abundance. We extend here that analysis to allow for intrinsic non-Gaussianities in the supernova probability distribution function, and discuss how this can be best modelled using the Bayes factor. Although it was shown that a precise measurement of σ<SUB>8</SUB> requires ∼10<SUP>5</SUP> SNeIa, current data already allow an important proof of principle. In particular, we make use of the 706 supernovae with z ≤ 0.9 of the recent Joint Lightcurve Analysis catalogue and show that a simple treatment of intrinsic non-Gaussianities with a couple of nuisance parameters is enough for our method to yield the values σ _8 = 0.84^{+0.28}_{-0.65} or σ<SUB>8</SUB> &lt; 1.45 at a 2σ confidence level. This result is consistent with mock simulations and it is also in agreement with independent measurements and presents the first ever measurement of σ<SUB>8</SUB> using SNeIa magnitudes alone.
false
[ "SNeIa magnitudes", "independent measurements", "non-Gaussianities", "Such signal", "SNeIa", "signal", "current data", "nuisance parameters", "principle", "structure", "z ≤", "the recent Joint Lightcurve Analysis catalogue", "noise", "a 2σ confidence level", "Bayes", "the supernova probability distribution function", "2σ", "first", "the cosmic microwave background", "Gaussianities" ]
11.228877
2.099466
-1
483093
[ "Ádámkovics, Máté", "Glassgold, Alfred E.", "Najita, Joan R." ]
2014ApJ...786..135A
[ "Shielding by Water and OH in FUV and X-Ray Irradiated Protoplanetary Disks" ]
29
[ "Astronomy Department, University of California, Berkeley, CA 94720, USA", "Astronomy Department, University of California, Berkeley, CA 94720, USA", "National Optical Astronomy Observatory, 950 North Cherry Avenue, Tucson, AZ 85719, USA" ]
[ "2014ApJ...788...66C", "2014ApJ...792....2D", "2015A&A...582A..88W", "2015ApJ...810..125G", "2015ApJ...810L..24S", "2016ApJ...817...82A", "2016ApJ...830..112B", "2016PhRvA..93f1401M", "2017A&A...602A.105H", "2017ApJ...836..242D", "2017ApJ...842...98D", "2017ApJ...847....6N", "2017ApJ...847...11W", "2018ApJ...855...98A", "2018ApJ...860..175W", "2018ApJ...862..122N", "2018JQSRT.204...12H", "2019ApJ...871..173A", "2019ApJ...874...90W", "2020ApJ...888..116S", "2020ApJ...890...93G", "2021AJ....161..217C", "2021AJ....162..185A", "2022A&A...665A..45P", "2022ApJ...930L..26B", "2023ARA&A..61..287O", "2023FaDi..245...52V", "2024A&A...682A..91V", "2024arXiv240614560T" ]
[ "astronomy" ]
8
[ "astrochemistry", "protoplanetary disks", "X-rays: stars", "Astrophysics - Earth and Planetary Astrophysics" ]
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[ "10.1088/0004-637X/786/2/135", "10.48550/arXiv.1403.8131" ]
1403
1403.8131_arXiv.txt
The gas and dust in protoplanetary disks are both important in the formation of stars and planets \citep[e.g.,][]{Dullemond2010, Williams2011}. Significant progress has been made in recent years by observing the gas with ground-based facilities and the {\it Spitzer} and {\it Herschel} space observatories. A variety of molecules have been detected in the inner planet forming regions of these disks, notably water and simple organic molecules \citep{Carr2004,CN08,CN11, Pascucci2009,Pontoppidan2010a,Pontoppidan2010b,Salyk2011,Najita2013}. The presence of water is of particular astrophysical and astrobiological significance. Water can be as abundant as CO in protoplanetary disks, making it the third most abundant molecule after H$_2$ and CO. Water and dust can also play important physical and chemical roles, a focus of this article. Numerous models of protoplanetary disk chemistry have been developed for the purpose of establishing and analyzing diagnostic spectral lines of the gas \citep{Markwick2002, Kamp2004, Nomura2005, Nomura2007, Nomura2009, Agundez2008, Gorti2008, Woods2009, Woitke2009a, Woitke2009b, Kamp2010, Fogel2011, Kamp2011, Thi2010,Heinzeller2011, Walsh2012, Aresu2011, Aresu2012, Bruderer2012, Meijerink2012, Akimkin2013}. These models tend to differ in significant physical details, e.g., how the gas is irradiated and how it is heated, as well as in the chemistry. In addition to our earlier paper \citep{GMN09}, several other models produce significant levels of water in the inner disk \citep[e.g.,][]{Woods2009, Thi2010, Fogel2011, Heinzeller2011, Aresu2011}, and thus support the idea of {\it in situ} formation in the gas phase. If this conclusion can be strengthened, as we attempt to do here, it would diminish the need for an origin of water from the transport of icy particles and bodies \citep[e.g.,\!\!~][]{Ciesla2006}. In 2009 Bethell and Bergin called attention to the possibility that the disk midplane is shielded from dissociating far ultraviolet (FUV) radiation by water and OH in the disk atmosphere \citep[henceforth BB09]{BB09}. They described how including such molecular shielding could account for several features of the water and OH emission observed by {\it Spitzer}. Molecular shielding becomes important when grains have settled out of the atmosphere and dust is no longer the dominant absorber of UV radiation. Shielding by molecules can limit the penetration of the FUV through the disk atmosphere and thus affect the transition from regions dominated by atomic oxygen to regions with oxygen-bearing molecules. To make these points BB09 adopted an FUV irradiated isothermal atmosphere and employed a simplified chemistry for OH and \water. In particular, they assumed hydrogen to be completely in molecular form. One potential difficulty with this assumption is that the formation of \water\ and OH requires H$_2$ as a precursor, and the formation of H$_2$ on grains may be limited when the grain surface area is reduced by grain settling. A more complete thermal-chemical model is, therefore, needed to explore how well shielding by \water\ and OH works when when grain settling is advanced. This topic has not been pursued by earlier thermal-chemical models. For example, the protoplanetary disk model (ProDiMo) of Woitke et al.\ (2009; see also Aresu et al. 2011, Meijerink et al. 2012), while detailed and inclusive of many thermal and chemical processes, includes molecular shielding by H$_2$ and CO, but not \water\ and OH. In this paper we incorporate UV irradiation and molecular shielding by \water\ and OH into our earlier thermal-chemical model of the inner disk atmosphere to explore role of \water\ and OH as UV opacity sources. We treat carefully a number of physical processes related to the abundance of water and OH. Unlike BB09, who prescribed two isothermal layers, we evaluate the temperatures in the transition between these two regions of the atmosphere by considering the chemistry-dependent thermal rate equations. In our previous study of water with a purely X-ray model \citep{GMN09}, we obtained significant levels of water in the inner regions of a typical T Tauri star disk, but much lower levels of OH than are found by {\it Spitzer} (NAG11) or in the model of BB09. Therefore, we also explore in this paper whether a simple UV-irradiated thermal-chemical model can account for the properties of both the \water\ and OH emission detected with {\it Spitzer}. In the next section we describe how FUV photodissociation is treated in our model, including how it adds to the heating of the disk atmosphere. One of our goals is to describe our assumptions and the underlying reasons for them in enough detail to allow useful comparison with the results of other investigations. In Section 2, we describe improvements to our model, most importantly the addition of UV photodissociation and its associated heating. We have also added photoelectric and H$_2$ formation heating to our model.
We have developed an integrated thermal-chemical model for protoplanetary disk atmospheres that includes irradiation by FUV and X-rays, grain settling, and a detailed treatment of the physical processes that affect water and other molecules. Our objective has been to identify some of the key processes and how they operate. We emphasized that several of the processes are uncertain, especially those relating to grains, and we have focused on a reference model that represents a typical CTTS disk. Our results support the idea that \water\ in the inner regions of protoplanetary disks can be formed {\it in situ}. This thesis has been discussed in several of the earlier modeling papers cited in Section 1. Especially noteworthy are those that obtain sufficient warm columns of \water\ to be consistent with the amounts deduced from Spitzer observations, e.g., Glassgold et al.~(2009), Woods \& Willacy (2009), Heinzeller et al.~2011, Najita et al. (2011). In the present work, both the physical and astrophysical details of how adequate amounts of \water\ can be formed are analyzed in detail. Among the new results presented here is the key role of dust surface area in regulating processes that directly affect the \water\ abundance, i.e., H$_2$ formation, dust-gas thermal accommodation, and the extinction of FUV radiation. All of these processes are sensitive to the size distribution of the grains in the upper atmosphere of the disk. We find that, in our updated model, large amounts of \water\ and OH are synthesized in the observable warm regions of a typical protoplanetary disk atmosphere even in the presence of FUV and X-ray radiation and substantial grain settling. The \water\ and OH are located primarily in distinct locations, with the warm \water\ layer located below a warmer layer of OH. Both the temperatures and total numbers of molecules are consistent with {\it Spitzer} observations of CTTS, with the OH in the atmosphere extending over a larger range of radii than the \water. In our model, freeze out on dust grains restricts the gaseous water reservoir in the warm atmosphere to radii $< 4$\,AU. UV absorption by \water\ and OH is important in accounting for the large amounts of warm water that are observed, confirming the ideas discussed by BB09. FUV radiation affects the width and sharpness of the atomic to molecular transition, whereas X-rays and grain settling affect the depth at which the transition starts. X-rays also generate ions that destroy water below the photodissociation region. Many of these conclusions can be tested with suitable observations, as has already been done in some cases. Together with the discussion in previous sections, they present a richer picture of an atomic to molecular in protoplanetary disk atmospheres. \vspace{2ex} We acknowledge support from NASA grant NNG06GF88G (Origins) and NASA grant 1367693 ({\it Herschel} DIGIT). We are particularly grateful to Paula D'Alessio for providing the disk model used in this work, and we will remember her for her kindness and her many contributions to the understanding of protoplanetary disks.
14
3
1403.8131
We present an integrated thermal-chemical model for the atmosphere of the inner region of a protoplanetary disk that includes irradiation by both far-ultraviolet (FUV) and X-ray radiation. We focus on how the photodissociation of H<SUB>2</SUB>O and OH affects the abundances of these and related species and how it contributes to the heating of the atmosphere. The dust in the atmosphere plays several important roles, primarily as the site of H<SUB>2</SUB> formation and by absorbing the FUV. Large amounts of water can be synthesized within the inner 4 AU of a disk around a typical classical T Tauri star. OH is found primarily at the top of a warm region where the gas temperature is T <SUB>g</SUB> ≈ 650-1000 K and H<SUB>2</SUB>O is found below it, where the temperature is lower, T <SUB>g</SUB> ≈ 250-650 K. The amounts of H<SUB>2</SUB>O and OH and the temperatures of the regions in which they formed are in agreement with recent Spitzer measurements and support the notion of the in situ production of water in the inner regions of protoplanetary disks. We find that the synthesized water is effective in shielding the disk mid-plane from stellar FUV radiation.
false
[ "stellar FUV radiation", "protoplanetary disks", "-", "g</SUB", "several important roles", "recent Spitzer measurements", "a typical classical T Tauri star", "SUB", "water", "plane", "situ", "irradiation", "the inner region", "the inner regions", "Tauri", "a protoplanetary disk", "Spitzer", "OH", "mid", "(FUV" ]
10.614265
12.406733
-1
745725
[ "Cooke, E. A.", "Hatch, N. A.", "Muldrew, S. I.", "Rigby, E. E.", "Kurk, J. D." ]
2014MNRAS.440.3262C
[ "A z = 2.5 protocluster associated with the radio galaxy MRC 2104-242: star formation and differing mass functions in dense environments" ]
60
[ "School of Physics and Astronomy, University of Nottingham, University Park, Nottingham NG7 2RD, UK", "School of Physics and Astronomy, University of Nottingham, University Park, Nottingham NG7 2RD, UK", "School of Physics and Astronomy, University of Nottingham, University Park, Nottingham NG7 2RD, UK", "Leiden Observatory, Leiden University, PO Box 9513, NL-2300 RA Leiden, the Netherlands", "Max-Planck-Institut für extraterrestrische Physik, Giessenbachstrasse, D-85748 Garching, Germany" ]
[ "2014ApJ...789...18K", "2015A&A...579A..89C", "2015A&A...582A..30P", "2015ApJ...802...31D", "2015ApJ...811..125T", "2015MNRAS.452.2318C", "2015MNRAS.452.2528M", "2016A&A...592A.161N", "2016A&A...596A.100P", "2016A&ARv..24...14O", "2016ApJ...816...83C", "2016ApJ...817..160L", "2016ApJ...826L..28H", "2016ApJ...833...15F", "2016ApJS..226...23Z", "2016MNRAS.456.3827O", "2016MNRAS.459..387H", "2016MNRAS.462..421H", "2017A&A...599A.123N", "2017ApJ...836..136F", "2017ApJ...845..172B", "2017MNRAS.464..876H", "2017MNRAS.470.2170K", "2018ApJ...862...64S", "2018MNRAS.473.1977S", "2018MNRAS.473.2335M", "2018MNRAS.474.4612L", "2018MNRAS.480.1340I", "2018PASJ...70S..32U", "2019ApJ...871...83S", "2019ApJ...879....9S", "2019JKAS...52..235S", "2019MNRAS.489.5381M", "2020ApJ...899....5I", "2020ApJ...899...79S", "2020MNRAS.496.3169A", "2021A&A...654A.121P", "2021AN....342..153D", "2021ApJ...907....3L", "2021ApJ...912...55S", "2021ApJ...912...60S", "2021ApJ...915...32S", "2021ApJ...916...35I", "2021ApJ...919...51M", "2021IAUS..359..126I", "2021MNRAS.502.4558C", "2021MNRAS.503.1507V", "2021MNRAS.504.5054A", "2022ApJ...930..102Y", "2022ApJ...933....9L", "2022ApJ...935...80L", "2022ApJ...941..134H", "2022ApJ...941..180W", "2022MNRAS.510..674W", "2022MNRAS.513.3252A", "2022Univ....8..554A", "2023MNRAS.518.4159D", "2024A&A...682A.186L", "2024ApJ...960..132L", "2024MNRAS.528.6934S" ]
[ "astronomy" ]
7
[ "galaxies: clusters: individual: MRC 2104-242", "galaxies: high-redshift", "Astrophysics - Astrophysics of Galaxies", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
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[ "10.1093/mnras/stu522", "10.48550/arXiv.1403.4259" ]
1403
1403.4259_arXiv.txt
Locally, the star formation rate (SFR)-mass relation does not change as a function of galaxy environment; the fraction of galaxies which are star forming differs but the specific star formation rate (sSFR) is constant irrespective of environment \citep{YPeng2010}. This SFR-mass relation evolves with redshift, however cluster and field galaxies continue to lie on the same relation up to $z=1$ \citep{Muzzin2012}. At higher redshifts, studies have found that this trend of a constant sSFR between galaxies in the process of forming a cluster (protocluster galaxies) and field galaxies appears to continue, implying a sSFR independent of environment \citep{Koyama2013a,Koyama2013b}. The existence of a ``main sequence" for galaxies suggests that star formation in galaxies proceeds in the same way in (proto)clusters as it does in the field, even at redshifts $z>2$. Protocluster galaxy properties, however, differ from those in the field: the progenitors of low redshift clusters have previously been found to contain member galaxies that are older, more star-forming, more metal-rich and twice as massive as field galaxies at the same redshift \citep{Steidel2005,Hatch2011b,Koyama2013a,Kulas2013}. This implies that cluster galaxies have experienced an accelerated growth in their early years, yet their sSFRs show no difference from the field up to redshift $z = 2$. Previously, the SFR-mass relation at $z>2$ has been studied using masses derived from K-band fluxes, and SFRs corrected using mass-dependent dust extinction estimates \citep{Koyama2013a,Koyama2013b}. Using a dust extinction law that is solely dependent on the mass of the object makes it difficult to find extreme starbursts that lie above the main sequence. Using the rest frame UV slope as a direct measure of dust extinction, as well as infrared star formation indicators such as 24\,\micron and 250\,\micron fluxes, may help to break this degeneracy between normal star-forming galaxies and heavily dust-obscured star-bursting objects. Combining this with SED-derived masses should provide a better measure of the SFR-mass relation for protocluster and field galaxies at $z >2$. In this paper we investigate the SFR-mass relation in a candidate protocluster field, around the radio galaxy MRC\,2104$-$242. This field was observed as part of an infrared survey of eight high-redshift radio galaxies (HzRGs), described in \citet{Galametz2010b} and \citet{Hatch2011a}. Four of these HzRGs appeared to be surrounded by an overdensity of red galaxies, one of which (MRC\,0156$-$252) has recently been spectroscopically confirmed to lie within a large-scale structure \citep{Galametz2013b}. Another of these targets, MRC\,2104$-$242, had a 3$\sigma$ overdensity % of red galaxies and the angular correlation function showed that the galaxies in this field were more clustered than average \citep{Hatch2011a}. MRC\,2104$-$242 lies at $z = 2.49$, which means the \Halpha emission line falls directly within the ISAAC narrow-band filter at 2.29\,$\mu$m. This allows us to select star-forming galaxies within a narrow redshift range ($\Delta z = 0.05$) around the radio galaxy. Using optical to MIR photometry we have studied the masses and star-forming properties of \Halpha selected galaxies around MRC\,2104$-$242. We have compared the results in the radio galaxy field to a control field sample, using the same selection techniques throughout. The outline of the paper is as follows: Section \ref{sec:data} outlines the observations, data reduction and sample selection. Section \ref{sec:properties} describes our methods in determining the galaxy properties. In Section \ref{sec:results} we present our results and look at galaxy properties as a function of environment. Section \ref{sec:discussion} discusses our key results and possible implications and Section \ref{sec:summary} presents a summary. We assume a $\Lambda$CDM cosmology with H$_0 = 70$\,km\,s$^{-1}$\,Mpc$^{-1}$, $\Omega_M = 0.3$ and $\Omega_\Lambda = 0.7$ throughout, unless stated otherwise. We adopt a \citet{Chabrier2003} initial mass function (IMF) for all our calculations and magnitudes are given in the AB system unless stated otherwise.
\label{sec:summary} We have undertaken a NB survey of the field around the HzRG MRC\,2104$-$242. We have selected star-forming galaxies in this field and compared their properties with those of a field sample at similar redshifts. Here we present our key results: \begin{enumerate} \parsep0pt \item The field around the HzRG MRC\,2104$-$242 is overdense compared to blank control fields, with a level of overdensity of $8.0\pm0.8$ times the average blank field, which is consistent with this field being the progenitor of a low redshift cluster, i.e. a protocluster. \item The protocluster galaxies around MRC\,2104$-$242 are more massive and have more hidden star formation than control field galaxies at the same redshift. When we take a mass selected field sample we find no difference in the SFR and sSFR between the two environments, and only a minor difference in the dust content. \item Star formation at $z \sim 2.5$ is governed predominantly by galaxy mass, not environment. After including dust-extincted star formation using 24\,\micron and \emph{Herschel} data we find that the average SFR-mass relations are the same irrespective of environment and both the protocluster and control field galaxies lie close to the main sequence. \item We find a large difference in the mass distributions between environments: we expect to find $\sim21$-$22$ galaxies in the protocluster at masses $M<10^{10}$\,M$_{\odot}$ and detect none. This could indicate a higher level of dust extinction in low mass galaxies in the protocluster. It may alternatively be due to galaxies in the protocluster forming more high mass galaxies through monolithic collapse or undergoing many more mergers in the early stages of their growth. \item We find tentative evidence of a larger fraction of starburst galaxies in the protocluster than in the control field. Further data is required to confirm the 250\,\micron detections, however a more rapid mode of star formation in denser environments may explain how protocluster galaxies build up their mass quicker than in the field. \item The overdensity we detect in this small area is highly dependent on the mass range we consider. It can range from an overdensity of $0$ (at $M < 10^{10}$\,M$_{\odot}$) to $55$ ($M > 10^{10.5}$\,M$_{\odot}$). It is important when quantifying protoclusters to compare their mass functions, rather than simply number overdensities. \end{enumerate}
14
3
1403.4259
We present results from a narrow-band survey of the field around the high-redshift radio galaxy MRC 2104-242. We have selected Hα emitters in a 7 arcmin<SUP>2</SUP> field and compared the measured number density with that of a field sample at similar redshift. We find that MRC 2104-242 lies in an overdensity of galaxies that is 8.0 ± 0.8 times the average density of a blank field, suggesting it resides in a large-scale structure that may eventually collapse to form a massive cluster. We find that there is more dust obscured star formation in the protocluster galaxies than in similarly selected control field galaxies and there is tentative evidence of a higher fraction of starbursting galaxies in the denser environment. However, on average we do not find a difference between the star formation rate (SFR)-mass relations of the protocluster and field galaxies and so conclude that the SFR of these galaxies at z ∼ 2.5 is governed predominantly by galaxy mass and not the host environment. We also find that the stellar mass distribution of the protocluster galaxies is skewed towards higher masses and there is a significant lack of galaxies at M &lt; 10<SUP>10</SUP> M<SUB>⊙</SUB> within our small field of view. Based on the level of overdensity we expect to find ∼22 star-forming galaxies below 10<SUP>10</SUP> M<SUB>⊙</SUB> in the protocluster and do not detect any. This lack of low-mass galaxies affects the level of overdensity which we detect. If we only consider high-mass (M &gt; 10<SUP>10.5</SUP> M<SUB>⊙</SUB>) galaxies, the density of the protocluster field increases to ∼55 times the control field density.
false
[ "galaxy mass", "galaxies", "galaxy", "starbursting galaxies", "similarly selected control field galaxies", "∼22 star-forming galaxies", "similar redshift", "M", "the protocluster galaxies", "low-mass galaxies", "the control field density", "higher masses", "the protocluster and field galaxies", "the protocluster field", "star formation", "these galaxies", "overdensity", "a field sample", "MRC", "a blank field" ]
13.528869
6.286981
159
12395976
[ "Steiner, A. W.", "Gandolfi, S.", "Fattoyev, F. J.", "Newton, W. G." ]
2015PhRvC..91a5804S
[ "Using neutron star observations to determine crust thicknesses, moments of inertia, and tidal deformabilities" ]
112
[ "Institute for Nuclear Theory, University of Washington, Seattle, Washington 98195, USA; Department of Physics and Astronomy, University of Tennessee, Knoxville, Tennessee 37996, USA; Physics Division, Oak Ridge National Laboratory, Oak Ridge, Tennessee 37831, USA", "Theoretical Division, Los Alamos National Laboratory, Los Alamos, New Mexico 87545, USA", "Department of Physics and Astronomy, Texas A&amp;M University-Commerce, Commerce, Texas 75429, USA; Department of Physics and Nuclear Theory Center, Indiana University, Bloomington, IN 47405, USA", "Department of Physics and Astronomy, Texas A&amp;M University-Commerce, Commerce, Texas 75429, USA" ]
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[ "astronomy", "physics" ]
6
[ "26.60.Kp", "21.65.Cd", "97.60.Jd", "Equations of state of neutron-star matter", "Asymmetric matter neutron matter", "Neutron stars", "Nuclear Theory", "Astrophysics - High Energy Astrophysical Phenomena", "Astrophysics - Solar and Stellar Astrophysics" ]
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[ "10.1103/PhysRevC.91.015804", "10.48550/arXiv.1403.7546" ]
1403
1403.7546_arXiv.txt
Recent neutron star (NS) mass and radius observations have provided new constraints on the neutron star mass-radius curve and on the equation of state (EOS) of dense matter~\cite{Lattimer01}. The EOS, in turn, is a fundamental property of quantum chromodynamics, which probes cold and dense matter which is otherwise difficult to access in experiment. In the near future, mass and radius observations may be complemented by other constraints on NS structure. Although thousands of pulsars have been observed, there is only one binary system where both NSs are radio-active pulsars, PSR J0737-3039. The ability to observe pulsations from both NSs and the extreme nature of the system~\cite{Burgay03,Kramer06}, enables a potential measurement of the moment of inertia of one of the neutron stars~\cite{Damour88}. Also, the Laser Interferometer Gravitational-Wave Observatory is expected to measure the gravitational wave signal from a NS merger within the near future~\cite{Abadie10}, and a sufficiently large signal to noise observation will enable a measurement of the neutron star tidal deformability~\cite{Flanagan08,Hinderer10,Damour12,delPozzo13, Read13,Lackey14} (denoted by $\lambda$ and sometimes also called ``tidal polarizability''). It turns out these two types of new observations are intimately related: the moment of inertia of a NS is strongly correlated with its tidal deformability~\cite{Yagi13a,Yagi13b}. NSs can accrete matter from main-sequence companions which results in the emission of X-rays and the heating of the NS crust. If the accretion stops (referred to as ``quiescence'' since the X-rays from accretion subside), then the cooling NS crust can be directly observed~\cite{Rutledge02}. The timescale for this cooling is proportional to the square of the NS crust thickness~\cite{Lattimer94}, and thus the crust thickness is important for determining the properties of the crust from observations of crust cooling~\cite{Lewin06,Shternin07,Brown09,Page13}. Another potential constraint of NS structure comes from pulsar glitches. Previous papers~\cite{Link99,Andersson12,Chamel13} have shown that, if NS crusts are believed to be the location of the angular momentum reservoir which contributes to the glitch spin up, then a significant fraction of the NS's moment of inertia must lie in the superfluid component of the crust. Thus glitches are sensitive to the crustal fraction of the moment of inertia, denoted $\Delta I/I$. Finally, many of these quantities are (at least weakly) correlated with the nuclear symmetry energy~\cite{Steiner05,Tsang12,Li14}. The nuclear symmetry energy is the difference between the energy per baryon of neutron matter and that of nuclear matter (we ignore quartic terms, see Ref.~\cite{Steiner06}). We denote the symmetry energy $S(n_B)$, where $n_B$ is the baryon number density, and $S\equiv S(n_0)$, where $n_0$ is the nuclear saturation density. The quantity $3 n_0 S^{\prime}(n_0)$ is denoted as $L$. The value of $L$ determines the pressure of neutron-rich matter at the saturation density. The pressure of neutron-rich matter, in turn, is related to all of the above NS structure quantities given above. For the first time, we use existing NS mass and radius observations to predict the expected ranges of NS properties such as moments of inertia, tidal polarizabilities and crustal thicknesses which are measurable by a diverse range of ongoing observational programs. We generate these expected ranges based on Monte Carlo simulations using parameterizations which explore the full variation which is possible given current uncertainties in the nature of dense matter. Our EOS models are based on recent progress in the microscopic calculation of neutron-rich matter near the nuclear saturation density. At high densities, we assume no additional correlation with matter at lower densities, and use models which allow for strong phase transitions. Our method is in contrast to several previous papers which have computed theoretical predictions of moments of inertia, crust thicknesses, and tidal polarizabilities for smaller samples of representative EOSs~\cite{Kalogera99,Lattimer01,Morrison04,Bejger05,Steiner05,Lattimer05, Lattimer07,Fattoyev10,Fattoyev13,Lattimer13}. Future observations will consitute direct tests of the theoretical framework we use and of the systematics of current mass-radius observations.
There is a quandary with pulsar glitches which originates in two results. The first is that some EOS models (such as that of Akmal-Pandharipande-Ravenhall~\cite{Akmal98}) have small enough crusts that the fraction of the NS's moment of inertia contained in the crust is somewhat small ($\Delta I/I<0.05$ for a 1.4-$\mathrm{M}_{\odot}$ NS). The second is that there is a large amount of entrainment of superfluid neutrons by the lattice~\cite{Chamel05,Chamel12}, thus the amount of matter in the crust which is not strongly coupled to the lattice is only 15\%$-$25\% of the total. Together, these limit the magnitude of pulsar glitches to be smaller than those already observed in the Vela pulsar which requires $\Delta I/I \geq 0.016$~\cite{Link99,Espinoza11,Andersson12,Chamel13}. As can be seen in the upper-right of Fig.~\ref{fig:hist2}, NS mass and radius data predict a similar outcome, and the quandary stands. If we assume, however, that systematic uncertainties invalidate current mass and radius observations (as implied by Ref.~\cite{Suleimanov11}) and use our third data set which only contains a measurement of the moment of inertia of $I=(70{\pm}7)~\mathrm{M}_{\odot}~\mathrm{km}^2$, then we find many models with $\Delta I/I > 0.09$ as also shown in the upper-right panel. Smaller mass NSs give even larger values of $\Delta I/I$. As with the tidal deformabilities above, assuming a measurement of $I=(90{\pm}9)~\mathrm{M}_{\odot}~\mathrm{km}^2$ implies that $\Delta I/I$ could be larger than 0.10. Values as large as 0.11 can be obtained for lower mass neutron stars. These large values of $\Delta I/I$ can accommodate the observations of glitches in Vela even with the most extreme amounts of entrainment obtained in Ref.~\cite{Chamel05}. A similar conclusion has also been obtained independently in Ref.~\cite{Piekarewicz14}. The thermal evolution of a NS crust as it cools depends on the hydrostatic structure of the crust (as well as on how photons and neutrons are transported). Frequently, crust cooling is studied by using a small subset of the full variation possible for the hydrostatic structure~\cite{Brown09,Page13}. We find that, even for a fixed NS mass and radius, there is still considerable variation (due to the uncertainty in the EOS of dense matter) in the thickness of the crust. This is shown in the lower right panel of Fig.~\ref{fig:hist2} where the probability distribution of the radius of a 1.4-$\mathrm{M}_{\odot}$ NS is plotted versus the crust thickness $\Delta R$. We find that, for a NS with an 11 km radius, the crust thickness varies by 42\%. This means that a more complete variation in the parameter space may be required to determine the properties of the crust from crust cooling observations. If a measurement of the moment of inertia of PSR J0737-3039A was far outside our predicted range, then that implies a conflict with the mass and radius observations. This conflict could be resolved with modification of strong-field GR. However, this modification may have to be finely tuned in order to modify the NS structure without spoiling the agreement with GR found in the observations of the post-Keplerian parameters in the PSR J0737-3039 system~\cite{Kramer09}. Current NS mass and radius observations are subject to several strong systematic uncertainties (as described in Refs.~\cite{Steiner10,Steiner13,Lattimer14b}) and a moment of inertia measurement outside our predicted range could shed light on these systematics. Our understanding of NS structure would be best served by several different kinds of observations with different systematic uncertainties so that no one effect could dominate the results. The same reasoning given above also holds true for measurements of tidal deformabilities, crust thicknesses, and crustal fractions of the moment of inertia. Also, there are other neutron star observations which we could have used, but these are unlikely to strongly modify our results. For example, there are constraints from pulse profile modeling on the neutron star PSR J0437-4715~\cite{Bogdanov13}, but these are more than likely consistent with our results so long as the mass of this particular star is near 1 M$_{\odot}$. In particular, the 68\% limit for the radius of a 1-$\mathrm{M}_{\odot}$ star from that measurement is 11.3 to 14 km, and this overlaps the ranges given for all of the models presented in Table I. The exception to this is if the systematic uncertainties in the qLMXB and PRE burst observations are so large that the associated constraints on mass and radius should be ignored {\em and} a moment of inertia measurement was made for a lower mass star which was relatively small (i.e. $I < 70~\mathrm{M}_{\odot}~\mathrm{km}^2$ for a 1.4 $\mathrm{M}_{\odot}$ neutron star. A large increase in the NS maximum mass, such as that implied by Refs.~\cite{vanKerkwijk11,Romani12}, would significantly change these results. Larger maximum masses imply larger radii (larger pressure is needed at smaller densities to compete with gravity as the mass becomes larger), and thus larger moments of inertia and tidal deformabilities.
14
3
1403.7546
We perform a systematic assessment of models for the equation of state (EOS) of dense matter in the context of recent neutron star mass and radius measurements to obtain a broad picture of the structure of neutron stars. We demonstrate that currently available neutron star mass and radius measurements provide strong constraints on moments of inertia, tidal deformabilities, and crust thicknesses. A measurement of the moment of inertia of PSR J0737-3039A with a 10% error, without any other information from observations, will constrain the EOS over a range of densities to within 50%-60%. We find tidal deformabilities between 0.6 and 6 ×10<SUP>36</SUP>g cm<SUP>2</SUP>s<SUP>2</SUP> (to 95% confidence) for M =1.4 M<SUB>⊙</SUB> , and any measurement which constrains this range will provide an important constraint on dense matter. The crustal fraction of the moment of inertia can be as large as 10% for M =1.4 M<SUB>⊙</SUB> permitting crusts to have a large enough moment of inertia reservoir to explain glitches in the Vela pulsar even with a large amount of superfluid entrainment. Finally, due to the uncertainty in the equation of state, there is at least a 40% variation in the thickness of the crust for a fixed mass and radius, which implies that future simulations of the cooling of a neutron star crust which has been heated by accretion will need to take this variation into account.
false
[ "crust thicknesses", "crusts", "neutron stars", "recent neutron star mass and radius measurements", "radius", "dense matter", "inertia reservoir", "superfluid entrainment", "a neutron star crust", "future simulations", "strong constraints", "inertia", "moments", "account", "accretion", "currently available neutron star mass and radius measurements", "1.4", "tidal deformabilities", "a large enough moment", "M =1.4 M<SUB>⊙</SUB" ]
4.790985
2.423152
-1
483158
[ "Herrero-Illana, Rubén", "Pérez-Torres, Miguel Á.", "Alonso-Herrero, Almudena", "Alberdi, Antxon", "Colina, Luis", "Efstathiou, Andreas", "Hernández-García, Lorena", "Miralles-Caballero, Daniel", "Väisänen, Petri", "Packham, Christopher C.", "Rajpaul, Vinesh", "Zijlstra, Albert A." ]
2014ApJ...786..156H
[ "A Multi-wavelength View of the Central Kiloparsec Region in the Luminous Infrared Galaxy NGC 1614" ]
19
[ "Instituto de Astrofísica de Andalucía-CSIC, P.O. Box 3004, E-18008 Granada, Spain", "Instituto de Astrofísica de Andalucía-CSIC, P.O. Box 3004, E-18008 Granada, Spain; Centro de Estudios de Física del Cosmos de Aragón (CEFCA), E-44001 Teruel, Spain; Dpto. de Física Teórica, Universidad de Zaragoza, E-50009 Zaragoza, Spain", "Instituto de Física de Cantabria, CSIC-Universidad de Cantabria, E-39005 Santander, Spain; Augusto González Linares Senior Research Fellow.", "Instituto de Astrofísica de Andalucía-CSIC, P.O. Box 3004, E-18008 Granada, Spain", "Centro de Astrobiología (INTA-CSIC), Ctra. de Torrejón a Ajalvir, km 4, E-28850 Torrejón de Ardoz, Madrid, Spain", "School of Sciencies, European University Cyprus, Diogenes Street, Engomi, 1516 Nicosia, Cyprus", "Instituto de Astrofísica de Andalucía-CSIC, P.O. Box 3004, E-18008 Granada, Spain", "Instituto de Física Teórica, Universidad Autónoma de Madrid, E-28049 Madrid, Spain", "South African Astronomical Observatory, P.O. Box 9, Observatory 7935 Cape Town, South Africa; Southern African Large Telescope, P.O. Box 9, Observatory 7935 Cape Town, South Africa", "Department of Physics and Astronomy, University of Texas at San Antonio, One UTSA Circle, San Antonio, TX 78249, USA", "Department of Physics, University of Oxford, Denys Wilkinson Building, Keble Road, Oxford OX1 3RH, UK", "Jodrell Bank Centre for Astrophysics, University of Manchester, Manchester M13 9PL, UK" ]
[ "2014PhDT.......331H", "2015A&A...580A..35G", "2015ApJ...799...11X", "2015MNRAS.454..766A", "2015MNRAS.454.3679P", "2016IAUS..315E..71S", "2016PASJ...68...20S", "2017ApJ...835..174S", "2017MNRAS.470L.112H", "2017MNRAS.471.1634H", "2018AJ....155...81S", "2018ApJ...868...10L", "2020MNRAS.494.5793K", "2021A&ARv..29....2P", "2021AcASn..62...68C", "2021ApJ...916...73S", "2022ApJS..261....3P", "2022ChA&A..46..330C", "2024MNRAS.529.4468L" ]
[ "astronomy" ]
8
[ "galaxies: individual: NGC 1614", "galaxies: nuclei", "galaxies: starburst", "infrared: galaxies", "radio continuum: galaxies", "supernovae: general", "Astrophysics - Astrophysics of Galaxies", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
[ "1938ApJ....88...52B", "1985ApJ...298L...7H", "1988A&A...203..278M", "1989ApJ...345..245C", "1989agna.book.....O", "1990AJ.....99.1088N", "1990ARA&A..28..215D", "1991ApJ...370..102C", "1991ApJ...383..164G", "1991rc3..book.....D", "1992ARA&A..30..575C", "1992MNRAS.259..709C", "1993ApJ...411..565C", "1994MNRAS.266..431P", "1995MNRAS.273..649E", "1995MNRAS.277..857P", "1996ApJ...468..191P", "1997ApJ...479..694V", "1998AJ....115.1693C", "1998ApJ...498..579G", "1999ApJS..123....3L", "2000A&A...357...13R", "2000ApJ...532..845A", "2000MNRAS.313..734E", "2001A&A...366..439K", "2001AJ....122.1213S", "2001ApJ...546..952A", "2001MNRAS.324..325M", "2002MNRAS.335L..23P", "2003A&A...410...53S", "2003AJ....125.1210A", "2003AJ....126.1607S", "2003MNRAS.344.1000B", "2005A&A...440..775K", "2005MNRAS.360.1055P", "2006AJ....131.1163S", "2006ApJ...638..938A", "2007ApJ...656..148A", "2007ApJ...659..314P", "2007ApJ...666..870C", "2008ApJ...685..211D", "2008ApJS..178..189W", "2009A&A...502..541E", "2009A&A...507L..17P", "2009ARA&A..47...63S", "2009MNRAS.399.1641P", "2010A&A...513A..11O", "2010A&A...514A...1I", "2010A&A...519L...5P", "2010ApJ...709..884Y", "2010ApJ...711..328D", "2011A&A...529A.106I", "2011A&A...535A..79H", "2011A&A...535A..93P", "2011A&A...536A..36A", "2011AJ....142...79M", "2011ApJ...740...95B", "2012A&A...539A.134B", "2012A&A...540A.116R", "2012ApJ...744....2A", "2012ApJS..203....9U", "2012MNRAS.420.2209V", "2013A&A...553A..72K", "2013A&A...556A..47H", "2013AJ....146...47I", "2013ApJ...765...78A", "2013MNRAS.432..998Z" ]
[ "10.1088/0004-637X/786/2/156", "10.48550/arXiv.1403.7201" ]
1403
1403.7201.txt
Luminous Infrared Galaxies (LIRGs; $10^{11} \leq L_\mathrm{IR} / L_\odot \leq 10^{12}$; $L_\mathrm{IR}=L[8\textrm{ --- }1000\;\mu\mathrm{m}]$) are known to closely follow the far-infrared (FIR) to radio correlation, and hence must host either a burst of star formation, or an active galactic nucleus (AGN) at their center, or both. Disentangling whether an AGN, or a starburst is the dominant heating mechanism is essential to understand the role of the IR-phase in galaxy evolution \citep[see, e.g.][]{alonso-herrero12,alonso-herrero13}. The central kpc regions of LIRGs are heavily enshrouded in dust, which prevents their study at optical wavelengths. Fortunately, at infrared wavelengths the extinction is significantly lower than in the optical, and radio is essentially extinction-free, which permits the study of the innermost regions of LIRGs, if the required spatial resolution is available. Also, sub-arcsecond imaging with \emph{Chandra} allows to image the inner central regions of LIRGs, thanks to the penetrating power of X-rays. Therefore, high-angular (sub-arcsecond) resolution observations of local ($D \lesssim 100$~Mpc) LIRGs at infrared, radio, and X-rays can be efficiently used to disentangle a putative AGN from a starburst. NGC\,1614 (IRAS\, 04315-0840) is a galaxy merger in a late stage of interaction, with strong tidal tails and only one obvious nucleus, though there is evidence for the remnant of a secondary one \citep{neff90, vaisanen12}. At a distance of 64\,Mpc (\citealt{devaucouleurs91}; 1\arcsec corresponds to 310\,pc), NGC\,1614 has an infrared luminosity $L_\mathrm{IR} \simeq 4\times10^{11} L_\odot$ \citep{sanders03}. Using the photometric information from the IRAS Bright Galaxy Sample \citep{sanders03} and from the NRAO VLA Sky Survey \citep{condon98}, the derived $q$-factor \citep{helou85} is 2.46, well within the FIR-to-radio correlation. The central kpc region of NGC\,1614 hosts a prominent circumnuclear ring of star formation of $\sim600$\,pc diameter, revealed in Pa$\alpha$ \citep{alonso-herrero01}. Recently, several authors have suggested that the ring is formed by an inner Lindblad resonance, where the gas is driven to it through the dust lanes \citep{olsson10,konig13}. The existence of an AGN at the center of NGC\,1614 is still a matter of debate: \citet{risaliti00} classified the hard X-ray emitting source in the central region of the galaxy as an AGN. However, the low signal-to-noise detection of the power-law continuum makes its interpretation uncertain \citep{olsson10}. Sub-millimeter array (SMA) observations of NGC\,1614 seem to indicate a nuclear, non-thermal component, but which cannot be ascribed solely to an AGN, or to a starburst (SB) \citep{wilson08}. \citet{yuan10} have classified NGC\,1614 as a starburst-AGN composite (albeit with a significantly larger contribution from the starburst), using their new optical classification scheme. More recently, \citet{vaisanen12} have used 3.3$\,\mu$m spatially resolved polycyclic aromatic feature (PAH) imaging and continuum diagnostics to argue that an obscured AGN can be ruled out, concluding that NGC\,1614 is a pure starburst. In this paper, we present sub-arcsecond angular resolution radio (3.6 and 6\,cm), mid-IR (8.7$\mu$m), optical (0.4 and 0.8$\mu$m), and \emph{Chandra} X-ray images of the central kpc region of NGC\,1614 to study in detail the central kpc region of this LIRG. Our main aim is to shed light on the AGN/SB controversy existing in the literature, as well as to discuss the striking morphological similarities between the radio and mid-IR images, which suggest a common origin for both emission mechanisms.
\subsection{Radio and infrared images} We show in Figure~\ref{fig:all} our 3.6\,cm continuum VLA image of NGC\,1614 from November 2004 and the 8.7\,$\mu$m continuum T-ReCS image at similar angular resolution. We also show the \emph{HST}/NICMOS continuum-subtracted Pa$\alpha$ image for comparison. The outer circle in the images covers essentially all of the emission at each wavelength, and has a radius of $\sim780$\,pc. We identify five regions within the circumnuclear ring: A, B, C, and D, which correspond to areas of strong emission, and N, which roughly delimits the nuclear region ($r \lesssim 90$\,pc). For convenience, Figure~\ref{fig:all} shows two additional regions: R, which corresponds to the whole ring, and T, which encompasses the entire region (ring and nucleus, R+N). We show the locus and size for each of those regions, as well as their integrated 3.6\,cm, 8.7\,$\mu$m, 3.3\,$\mu$m and Pa$\alpha$ fluxes, in Table \ref{table:regionsandfluxes}. The most conspicuous feature is the prominent mid-IR emission of the nucleus, N, which contrasts with its rather faint emission at radio wavelengths (see Fig.~\ref{fig:all} and Table~\ref{table:regionsandfluxes}). The regions to the northwest of the ring, A and B, show a mid-IR/radio ratio below the average of the whole ring (R), while regions C and D show the opposite behavior. Figure~\ref{fig:azimuthal} shows the azimuthal profiles at all three wavelengths, starting from the central pixel (the brightest pixel in region N for the 3.6\,cm and $8.4\,\mu$m images) and towards eight cardinal directions, separated by 45\degs\ from each other. To adequately compare the profiles, we normalized the radio continuum and Pa$\alpha$ fluxes to the median of the ratios (8.7\,$\mu$m/3.6\,cm) of each region for the case of the mid-IR image, and to the median of the ratios (Pa$\alpha$/3.6\,cm) for the Pa$\alpha$ image. We therefore increased the values of the 3.6\,cm and Pa$\alpha$ values by factors of $19.26$ and $10.79$, respectively, which allows to see more clearly variations in the whole circumnuclear region. Figure~\ref{fig:azimuthal} also shows that both the continuum 8.7\,$\mu$m and Pa$\alpha$ emission follow almost exactly the same trend from the very center up to the outermost regions of the star-forming ring, as found by \citet{diaz-santos08}. Overall, the same trend is also seen when comparing the IR azimuthal profiles against those of the radio continuum. The remarkable morphological similarities seen at both radio and mid-infrared wavelengths in most regions of the circumnuclear ring strongly suggest that the mechanisms responsible for that emission must be related. \citet{konig13} show a similar plot (see their Fig.~6), with the azimuthal profile of the CO\,(2--1) and Pa$\alpha$ emission displayed together with the radio emission at three different bands. The peaks of radio and mid-infrared emission in the ring would pinpoint then the regions where most of the starburst activity has taken place in the last 10-20 Myr. The 8.7\,$\mu$m mid-IR flux includes both warm dust continuum and PAH emission. The PAH emission is known to vary quite significantly from one SF region to another, or one galaxy to another, depending on the exact physical conditions \citep[see, e.g.,][]{calzetti07, diaz-santos10}. There are morphological differences between the $3.3\mu$m PAH feature map and the rest of the IR and radio images. \citet{imanishi13} have suggested that either an age differentiation, or different dust extinction between the regions may explain those differences. \citet{olsson10} ascribed the radio emission at 3.6 and 6.0\,cm from NGC\,1614 mainly to free-free emission from H\,{\sc ii} regions (mainly due to massive stars). As we shall show in the next sections, the non-thermal contribution from core-collapse supernovae (CCSNe) and supernova remnants (SNRs) is also significant within the circumnuclear ring, implying the existence of young starbursts in the ring of NGC\,1614. \begin{figure*}\centering %\epsscale{.80} \includegraphics[width=0.9\textwidth]{azimuthalnew2.eps} \caption{Azimuthal profile of the fluxes in mid-IR (8.7\,$\mu$m, in blue), radio (3.6\,cm from Nov 2004, in green), Pa-$\alpha$ (1.9\,$\mu$m, in red) and continuum-subtracted PAH feature (3.3\,$\mu$m, in cyan) , starting from the center of the images. Radio, Pa-$\alpha$ and PAH fluxes are scaled up by the median of the 8.7\,$\mu$m/3.6\,cm, 8.7\,$\mu$m/Pa$\alpha$ and 8.7\,$\mu$m/3.3\,$\mu$m ratios, respectively (factors 19.2, 10.8 and 13.6). The profile is shown for eight cardinal directions, with an azimuthal binning of the size of a pixel (0.116\asec for the $3.3\,\mu$m image and 0.089\asec for the other cases). The shaded area, from 0.3 to 1.0\,arcsec, corresponds to the approximate width of the star formation ring.} \label{fig:azimuthal} \end{figure*} \begin{deluxetable*}{rrrrrrrrrr} \tabletypesize{\scriptsize} \tablecaption{Flux density and spectral index for the (circum)-nuclear region of NGC\,1614} \tablewidth{0pt} \tablehead{ %\colhead{} &\multicolumn{8}{c}{Integrated fluxes} \\ %\cline{2-9} \\ \colhead{} & \multicolumn{2}{c}{6.0\,cm} & \colhead{} & \multicolumn{3}{c}{3.6\,cm} & \colhead{} & \multicolumn{2}{c}{Spectral index} \\ \cline{2-3}\cline{5-7}\cline{9-10} \\ \colhead{} & \colhead{1986} & \colhead{1999} & \colhead{} & \colhead{1999} & \colhead{2004} & \colhead{2006} & \colhead{} & \colhead{Total} & \colhead{Non-thermal} \\ \colhead{Region} & \colhead{(mJy)} & \colhead{(mJy)} & \colhead{} & \colhead{(mJy)} & \colhead{(mJy)} & \colhead{(mJy)} & \colhead{} & \colhead{average} & \colhead{average} } \startdata A & $ 1.40 \pm 0.06 $ & $ 1.46 \pm 0.05 $ & & $ 1.29 \pm 0.05 $ &$ 1.26 \pm 0.04 $ & $ 1.23 \pm 0.04 $ & & $-0.48\pm0.06 $ & $-0.56\pm0.10$ \\ B & $ 3.42 \pm 0.12 $ & $ 3.71 \pm 0.11 $ & & $ 3.09 \pm 0.10 $ &$ 3.01 \pm 0.09 $ & $ 2.88 \pm 0.09 $ & & $-0.53\pm0.17 $ & $-0.72\pm0.16$ \\ C & $ 1.15 \pm 0.05 $ & $ 1.23 \pm 0.04 $ & & $ 0.91 \pm 0.04 $ &$ 1.00 \pm 0.03 $ & $ 0.97 \pm 0.03 $ & & $-0.74\pm0.04 $ & $-1.44\pm0.13$ \\ D & $ 1.22 \pm 0.06 $ & $ 1.13 \pm 0.04 $ & & $ 0.98 \pm 0.04 $ &$ 1.08 \pm 0.04 $ & $ 1.00 \pm 0.04 $ & & $-0.62\pm0.03 $ & $-1.14\pm0.22$ \\ N & $ 2.57 \pm 0.10 $ & $ 2.81 \pm 0.09 $ & & $ 1.85 \pm 0.07 $ &$ 1.79 \pm 0.06 $ & $ 1.79 \pm 0.06 $ & & $-0.68\pm0.13 $ & $-1.30\pm0.53$ \\ T & $ 34.06 \pm 1.06 $ & $ 35.70 \pm 1.08 $ & & $ 25.61 \pm 0.79 $ &$ 26.39 \pm 0.80 $ & $ 25.45 \pm 0.77 $ & & $-0.73\pm0.25 $ & $-1.19\pm0.63$ \\ R & $ 31.48 \pm 0.98 $ & $ 32.89 \pm 0.99 $ & & $ 23.77 \pm 0.73 $ &$ 24.60 \pm 0.74 $ & $ 23.66 \pm 0.72 $ & & $-0.73\pm0.25 $ & $-1.18\pm0.63$ \enddata \label{table:variability} \tablecomments{ Both images at 6.0\,cm were mapped with a beam of $0.71\times0.43$\,arcsec and the three epochs at 3.6\,cm with a beam of $0.50\times0.36$\,arcsec. The (two-point) spectral index was obtained for the 6.0 and 3.6\,cm simultaneous observations of NGC\,1614 carried out in 1999. Col. 7 shows the average spectral indices obtained directly from the spectral index map, while col. 8 spectral indices are obtained from the isolated non-thermal emission (see main text for details).} \end{deluxetable*} \subsection{Radio variability and spectral index of the circumnuclear ring in NGC\,1614} In star forming regions, non-thermal radio emission and variability are usually good tracers of recently exploded supernovae. We show in Fig.~\ref{fig:variability} the radio interferometric images of NGC\,1614 at 6.0\,cm (epochs 1986 and 1999) and 3.6\,cm (epochs 1999, 2004 and 2006), obtained with the VLA in A configuration, and imaged using the same restoring beam at all epochs, for consistency. In Table~\ref{table:variability}, we show the integrated flux for each defined region and for each epoch. \begin{figure*} %\epsscale{.80} \includegraphics[width=\textwidth]{variabilitynewlabel.eps} \caption{Sub-arcsecond resolution radio continuum images of NGC\,1614. Top panels show images at a wavelength of 6\,cm from observations in 1986 (left) and 1999 (right). Bottom pannels show 3.6\,cm maps from observations in 1999 (left; quasi-simultaneously taken also at 6\,cm), 2004 (center) and 2006 (right). Images of the same frequency were mapped with the same beam. In all cases, the bulk of the radio emission is within a circumnuclear star-forming ring of radius $\sim390$\,pc. Note that the peaks of the brightest regions at 3.6\,cm have a maximum whose position coincides rather well with the peaks seen at 6\,cm, but for region B, whose maxima clearly peak at different positions. The color scale is the same for each epoch but independent for each wavelength, corresponding to [0.1, 5.0]\,mJy/beam for the 6\,cm band and [-0.17, 2.65]\,mJy/beam for 3.6\,cm. See main text and Table \ref{table:variability} for details.} \label{fig:variability} \end{figure*} The images at each wavelength look overall similar, and the total radio emission at 3.6 and 6.0 cm has not varied among the epochs, within the uncertainties. In fact, while there seems to be an apparent increase of the flux density between 1986 and 1999 at 6\,cm (see Fig.~\ref{fig:variability}, top panel) for several regions, including N and B (the brightest spot in the circumnuclear ring), this variability is not quantitatively significant ($\lesssim1\sigma$, see Table~\ref{table:variability}) and hence we cannot claim they are real. Similarly, the images at 3.6\,cm between 1996 and 2006 (Fig.~\ref{fig:variability}, bottom panel) suggest that the flux density in region B has been steadily decreasing from 1996 till 2006, while region C would have experienced a rise and decrease of flux density during this period, possibly indicating supernova activity. Again, the variations are not significant, and unfortunately we lack further observations that could have allowed us to confirm, or rule out, those variations. We note, however, that the peaks of the regions sometimes show significant changes from epoch to epoch, suggesting that supernova events may be occurring. Still, it seems that no very radio bright, Type IIn supernova ($L_{\nu, \mathrm{peak}} \gsim 10^{28}$\ergshz) has exploded in NGC\,1614 during the 1986-2006 period. Such bright supernovae are known to evolve slowly and stay bright for $\gsim$10 yr, e.g. SN~1986J in NGC\,891 \citep{perez-torres02a}, or some of the supernovae in the compact nuclear starbursts of Arp~299-A \citep{perez-torres09b, bondi12} and Arp 220 \citep{parra07, batejat11}. Yet, we cannot exclude completely the possibility of a Type IIn SN having exploded and be missed by us, given the scarcity of the radio observations. We used the quasi-simultaneous 3.6 and 6.0\,cm observations in 1999 to derive the average spectral index $\alpha$ ($S_\nu\propto\nu^\alpha$) for each of the regions defined in Figs.~\ref{fig:all} and \ref{fig:variability}. The values in col.~7 of Table~\ref{table:variability} result from obtaining an average value of each region from the spectral index map created with the actual observations at 6.0 and 3.6\,cm from 1999 (which are the combination of the thermal free-free and non-thermal synchrotron radio emission), while col.~8 shows the intrinsic synchrotron radio spectral index, once the thermal component is subtracted from the total radio emission (see section 3.5 for details). The overall non-thermal spectral index for both the ring and the total area is $\alpha \simeq -1.2$. Such steep spectral index is suggestive of most of the diffuse, extended synchrotron radio emission being due to supernovae and supernova remnants. Note, however, that the brightest regions in the ring have spectral indices significantly flatter than $-1.2$, which might suggest that the synchrotron radio emission is mostly powered by SN remnants, rather than by recently exploded SNe. The [Fe\,{\sc ii}] emission line at 1.26 and 1.64\,$\mu$m has also been proposed as a tracer of the supernova rate in nearby starburst galaxies \citep[see, e.g.,][]{moorwood88, greenhouse91, colina92, colina93, vanzi97, alonso-herrero03}, both in a pixel-by-pixel basis, and as an integrated approach. The [Fe\,{\sc ii}] 1.26\,$\mu$m emission in NGC\,1614 shows a C-shape morphology in the circumnuclear star-forming ring \citep{rosenberg12}, with the brightest emitting region matching approximately region D in our 3.6 and 6.0\,cm images, i.e., an apparent anti-correlation between the maxima in the continuum VLA radio emission and the [Fe\,{\sc ii}] 1.26\,$\mu$m. However, after correcting for extinction the [Fe\,{\sc ii}] image, this anti-correlation disappears (Rosenberg, private communication). \subsection{Spatially resolved X-ray emission}\label{sec:xray} We fitted the \emph{Chandra} spectrum using standard procedures within the X-ray software XSPEC\footnote{http://heasarc.gsfc.nasa.gov/xanadu/xspec/} version 12.7.0. We fitted the data using four different models: (i) a pure thermal model (MEKAL), where the thermal emission is responsible for the bulk of the X-ray energy distribution; (ii) an absorbed power-law model (PL), which corresponds to a non-thermal source representing an AGN; (iii) a composite of a thermal plus an absorbed power-law model (MEPL); and (iv) a thermal model with tuned individual abundances (VMEKAL), which models the metal abundance pattern of type II SNe \citep[see, e.g.,][]{iwasawa11, zaragoza-cardiel13}. In all models, we kept fixed the Galactic absorption to the predicted value using the {\sc nh} tool within {\sc ftools} \citep{dickley90, kalberla05}. Neither the MEKAL nor the PL models yielded satisfactory fits to the data. The MEPL model fitted the data well, but with a physically unrealistic power-law index of $\Gamma=3.1$. Using a MEPL model to fit \emph{XMM-Newton} data (with a circular aperture of $\sim15\asec$), \citet{pereira-santaella11} obtained a luminosity four times higher, with $\Gamma=2$. Our best fit turned out to be the VMEKAL, which is shown in Fig.~\ref{fig:chandraall}, together with the soft and hard X-ray maps. This model finds abundances for Mg\,{\sc xi} (1.36\,keV), Si\,{\sc xiii} (1.85\,keV) and S\,{\sc xv} (2.4\,keV). The corresponding luminosity in the soft band for this model is $\log L\mathrm{(0.5-2 keV)}=40.78^{+0.04}_{-0.05}$\,erg/s, while there are not enough counts in the hard band to derive a realistic value for the luminosity. The fitted temperature is $kT=1.7\pm0.7$\,keV. \begin{figure*} \includegraphics[width=\textwidth]{chandra_cont_spec4.eps} \caption{\emph{Chandra} map (left) and spectrum fit (right) for NGC\,1614. The map shows the image of the total X-ray emission with the overlapped contours of the soft band (0.5--2.0\,keV, in green) and the hard band (2.0--10.0\,keV, in blue). The black circle corresponds to region T. Note that the emission is significantly more compact in the hard band than in the soft band.} \label{fig:chandraall} \end{figure*} \subsection{Thermal free-free and non-thermal (synchrotron) radio emission}\label{sec:radioemission} The bulk of the continuum radio emission observed in the central region of NGC\,1614 comes from its circumnuclear ring (Fig.~\ref{fig:all}), where a strong burst of star-formation is ongoing \citep{alonso-herrero01}. Massive stars and their associated H\,{\sc ii} regions would be responsible for the thermal free-free radio emission, while supernovae and supernova remnants would account for the non-thermal synchrotron radio emission. Disentangling the contribution from each of those two components is non-trivial from radio measurements alone, since it would require observations at several frequencies in the range from $\sim20$\,cm down to $\lsim$1\,cm, and ideally with the same angular resolution, which is not our case. Here, we estimate the expected thermal free-free continuum radio emission from extinction corrected Pa$\alpha$ measurements. Then, using our (extinction-free) continuum radio data at 3.6\,cm from November 2004, we infer the amount of radio emission that is of non-thermal, synchrotron origin. As a bonus, from the Pa$\alpha$ measurements we obtain one of the most relevant physical parameters in the starburst in NGC\,1614, namely its Lyman photon ionizing flux, $N_\mathrm{ion}$, which is used in Section~\ref{sec:sedfit} to compare with the value derived from the SED fitting of the starburst in the circumnuclear region of NGC\,1614. In fact, using standard relations \citep[see ][]{colina91} and the Pa$\alpha$ to H$\alpha$ recombination ratio and assuming no photon leakage, we can derive the ionizing photon flux, $N_\mathrm{ion}$, as: \begin{equation} N_\mathrm{ion} = 6.27\times10^{12}L_{\mathrm{Pa\alpha}} \ {\rm s^{-1}}, \end{equation} where $N_\mathrm{ion}$ is measured in photons/s and $L_{\mathrm{Pa\alpha}}$ is the Pa$\alpha$ extinction-corrected luminosity, in erg/s. From our continuum-subtracted Pa$\alpha$ image, we obtain a flux density of $\sim 47.1$\,mJy for the emission of the whole region, T, which corresponds to an absorbed $L_{\rm Pa\alpha} = 3.6\times10^{41}\ergs$. This luminosity translates into an (absorbed) ionizing photon flux of $N_\mathrm{ion} \approx 2.27\EE{54} \ {\rm s^{-1}}$. To obtain the relevant, unabsorbed ionizing photon flux, we corrected for the extinction, $A_V$. Fortunately, the extinction towards NGC\,1614 is well studied, and is in the range $A_V = 3 - 5$ \citep[see, e.g.,][]{neff90, puxley94, alonso-herrero01, kotilainen01, rosenberg12}. Assuming a value of $A_V=4$, and using our observed Pa$\alpha$, we obtain the unabsorbed Pa$\alpha$ flux by using standard H\,{\sc i} recombination lines ratios, for Case B \citep{baker38}, and using that \begin{equation} \frac{F(\lambda)}{I(\lambda)} = 10^{-C(H\beta)\left[f(\lambda)+1\right]}, \end{equation} where $F(\lambda)$ and $I(\lambda)$ are the absorbed and unabsorbed fluxes, respectively, $C(H\beta)$ is the reddening coefficient, and $f(\lambda)$ is the reddening function \citep{cardelli89}. The resulting unabsorbed Pa$\alpha$ flux is $\simeq80.2$ mJy, corresponding to an unabsorbed ionizing photon flux of $N_\mathrm{ion} \approx 3.87\EE{54} \ {\rm s^{-1}}$. Using standard relations between Pa$\alpha$ and H$\beta$ \citep[e.g.,][]{osterbrock89} and Eq.~3 from \citet{condon92}, we can obtain the thermal continuum radio emission as: \begin{equation} S_\mathrm{thermal} = 1.076\EE{13} \times F(\mathrm{Pa}\alpha) \,\nu^{-0.1}, \end{equation} with $S_\mathrm{thermal}$ in mJy, the unabsorbed Pa$\alpha$ flux, $F(\mathrm{Pa}\alpha)$, in erg\,cm$^{-2}$\,s$^{-1}$, and $\nu$ in GHz, and where we have assumed for simplicity a temperature of $10\,000$\,K and $N_\mathrm{e} = 10^4$\,cm$^{-3}$, which are typical of compact (i.e., size $\lesssim 1$\,pc) starburst regions. (The uncertainty in our estimates is dominated by the plasma electron temperature, since $S_{\rm th} \propto T_e^{0.52} $. A value of $T_e = 20000$ K would result in a thermal continuum radio flux $\sim 23$\% higher.) In this way, we isolated the thermal and non-thermal contributions to the radio emission, which are shown in Table~\ref{table:radioflux} and in Fig.~\ref{fig:thermalfraction}. The corresponding thermal free-free radio flux density is 11.03\,mJy at 3.6\,cm for the whole region, T, and is about 42\% of the total 3.6\,cm radio emission in the central regions of NGC\,1614. Since our Pa\,$\alpha$ measurements are much less affected by extinction than optical measurements, our decomposed values have a much weaker dependence on the actual value of the extinction. Indeed, allowing for an extinction $A_V$ in the range $(3-5)$, results in thermal free-free radio flux densities in the range $(9.65-12.60)$\,mJy. In summary, we obtained the thermal radio emission from a scaled version of the Pa$\alpha$ image, and the non-thermal radio emission as the result of the subtraction of the thermal emission from the total radio emission at 3.6\,cm from November 2004. The ratio of thermal free-free to synchrotron radio emission can be used as an indicator of the starburst age of each region in the circumnuclear ring of NGC\,1614. Regions where there is essentially no synchrotron radio emission imply that supernovae have not yet started to explode, indicating ages of at most $\sim4$\,Myr, while regions where the supernovae have already started to explode would be older. The models of P\'erez-Olea \& Colina (1995) provide a quantitative estimate of the thermal free-free emission from massive stars, and non-thermal radio continuum emission from supernovae and supernova remnants. The ratios of thermal to non-thermal radio emission (see Table~\ref{table:radioflux} and Fig.~\ref{fig:thermalfraction}) for regions A and B are about $\sim0.5$, while those of regions C and D of 1.1 and 1.2, respectively. The ratios above (and the free-free thermal continuum luminosities) can be well explained if the emission in regions C and D come from instantaneous bursts with ages $\lsim$5.5 Myr, where supernovae have only recently started to explode. On the other hand, the emission from regions A and B would come from slightly older ($\sim8$\,Myr) bursts, where essentially all exploding supernovae come from stars with masses in the 20-30 \msun\ range. We note that the above discussion is valid at 3.6\,cm and is made under the assumption of a constant extinction of $A_V=4$ and variations across the ring may affect the thermal to non-thermal ratios. In fact, there seems to exist a gradient of the extinction increasing towards the west, as shown in Fig.~1 of \citet{konig13}. Considering an extinction of $A_V=5$ for regions A and B, and $A_V=3$ for regions C and D, we still obtain thermal to non-thermal radio emission ratios of $\sim0.6$ for regions A and B, $\sim0.7$ for region C, and $\sim0.8$ for region D. The images of the decomposed radio emission help to understand the apparent paradox of the prominent mid-IR emission of the nucleus, N, which shows rather faint emission at radio wavelengths (see Fig.~\ref{fig:all} and Table~\ref{table:regionsandfluxes}). Since the thermal free-free emission is directly proportional to the Pa-$\alpha$ flux, the decomposed images would indicate that the radio emission from the nuclear region is dominated by thermal free-free emission (for any $A_V$ in the range $3-5$), which in turn suggests it is powered by a starburst, rather than an AGN, as we show in Section~\ref{sec:agn}. In the absence of low-frequency absorption, we would expect that the emission at lower frequencies (e.g., 21\,cm) should be dominated by synchrotron emission, in contrast to the 3.6 and 6.0\,cm images, where the contribution of the thermal emission is relevant. In fact, when we scale the non-thermal emission at 3.6\,cm (Fig.~\ref{fig:thermalfraction}) to 21\,cm, using our derived non-thermal spectral indices (see Table~\ref{table:variability}), one would expect to recover the emission from the MERLIN image at 21\,cm shown in Fig.~4 in \citet{olsson10}. Although the extrapolated image correlates, in general terms, well with the Olsson et al. image, all regions of NGC\,1614, except regions A and N, show at 21\,cm a lower flux density than expected, with the radio emission from region D being especially suppressed. This should not come as a surprise, as those starburst regions have many massive stars that create big H\,{\sc ii} regions around them. Those H\,{\sc ii} regions are very efficient low-frequency absorbers, as demonstrated by, e.g., the large emission measure (EM) values in the vicinities of SN2000ft in the circumnuclear starburst of NGC 7469 \citep{alberdi06, perez-torres09a}, or around supernova A0 in Arp 299A \citep{perez-torres09b}. Indeed, the low-frequency absorption implies that region D, the one showing the least flux density at 21\,cm, has an EM $\approx 1.2 \times 10^7$ cm$^{-6}$\,pc, a value very similar to that found for in the vicinities of supernovae SN 2000ft in NGC7469, or A0 in Arp 299A. Regions B and C have moderate EM values ($\approx [3.3, 4.7] \times 10^6$ cm$^{-6}$\,pc). Finally, regions A and N have negligible EM values and essentially their radio emission is not being efficiently suppressed at low frequencies. While a discussion of the specific reasons for the differences in the low-frequency absorption displayed by the circumnuclear regions of NGC 1614 is beyond the scope of this paper, we just note here that our results are in agreement with all regions being synchrotron dominated at 21\,cm, but whose emission is being significantly suppressed in some regions by low-frequency absorption that is most likely due to foreground absorbers, i.e., H\,{\sc ii} regions. The main exceptions are regions A and N, which seems to suffer very little low-frequency absorption, possibly due to a smaller density in those region, as indicated by the low EM values. \begin{figure*} %\epsscale{.80} \includegraphics[width=\textwidth]{thermalfraction.eps} \caption{Decomposition of the 3.6\,cm flux from November 2004 of the central region of NGC\,1614 into thermal (left panel, scaled version of the Pa$\alpha$ applying an uniform extinction of $A_V=4$) and non-thermal (middle panel, thermal radio emission subtracted from the total radio flux) components. The right panel shows the relative contribution of the thermal emission. Note that regions A and B are dominated by synchrotron non-thermal emission, in contrast with regions C and D.} \label{fig:thermalfraction} \end{figure*} \begin{deluxetable*}{rrrrrrr} \tabletypesize{\scriptsize} \tablecaption{Thermal and non-thermal radio emission in NGC\,1614} \tablewidth{0pt} \tablehead{ \colhead{} & \colhead{$S_\mathrm{th}$} & \colhead{$S_\mathrm{syn}$}& \colhead{$L_\mathrm{th}$}& \colhead{$L_\mathrm{syn}$}& \colhead{Unabs. $L_{\mathrm{Pa}\alpha}$}& \colhead{$N_\mathrm{ion}$}\\ \colhead{Region} & \colhead{(mJy)} & \colhead{(mJy)} & \colhead{($10^{27}$ erg s$^{-1}$ Hz$^{-1}$)} & \colhead{($10^{27}$ erg s$^{-1}$ Hz$^{-1}$)} & \colhead{($10^{40}$ erg s$^{-1}$ Hz$^{-1}$)}& \colhead{($10^{53}$ s$^{-1}$)} \\ \colhead{(1)} & \colhead{(2)}& \colhead{(3)} & \colhead{(4)} & \colhead{(5)} & \colhead{(6)}& \colhead{(7)} } \startdata A & 0.45 & 0.96 & 2.22 & 4.69 & 2.55 & 1.60 \\ B & 0.92 & 2.03 & 4.49 & 9.97 & 5.15 & 3.23 \\ C & 0.51 & 0.46 & 2.49 & 2.24 & 2.86 & 1.79 \\ D & 0.58 & 0.48 & 2.86 & 2.37 & 3.29 & 2.06 \\ N & 1.01 & 0.69 & 4.97 & 3.38 & 5.70 & 3.58 \\ T & 11.03 & 15.47 & 54.03 & 75.79 & 62.07 & 38.92 \\ R & 10.01 & 14.77 & 49.07 & 72.41 & 56.37 & 35.34 \enddata \label{table:radioflux} \tablecomments{Col. 1: Region name; Col. 2: Thermal fraction of the flux at 3.6\,cm, obtained as a scaled version of the Pa$\alpha$ image assuming a constant extinction of $A_V=4$); Col. 3: Synchrotron fraction of the flux at 3.6\,cm, (obtained as Col.\,1 subtracted from the total radio emission); Col.4: Thermal radio luminosity at 3.6\,cm; Col. 5: Synchrotron radio luminosity at 3.6\,cm; Col. 6: Unabsorbed Pa$\alpha$ luminosity; Col. 7: Number of ionizing photons.} \end{deluxetable*} %%% SED fit to the SB in NGC\,1614 \begin{figure}\centering %\epsscale{.80} \includegraphics[width=\columnwidth]{NGC1614.ps} \caption{NGC\,1614 SED fitting. Photometric data points are plotted in blue, while \emph{Spitzer} IRS spectrum is shown in pink. Lines show the overall fit (solid black), the starburst contribution (dashed red) and the host galaxy contribution (dot-dashed green).} \label{fig:andreas} \end{figure} \begin{deluxetable}{rr} \tabletypesize{\scriptsize} \tablecaption{SED Model fitting} \tablewidth{0pt} \tablehead{ \colhead{Parameter} & \colhead{Value}} \startdata $L_\mathrm{SB}$ & $10^{11.39}L_\odot$ \\ e-folding time of SB & $35.4$\,Myr \\ Age of SB & $29.5$\,Myr \\ $\mathrm{SFR_\mathrm{max}}$ & $85.1 M_\odot\,\mathrm{yr}^{-1}$ \\ $\mathrm{SFR_\mathrm{mean}}$ & $57.7 M_\odot\,\mathrm{yr}^{-1}$ \\ & (averaged over 29.5\,Myr) \\ $\nu_\mathrm{SN}$ & $0.43\,\mathrm{yr}^{-1}$ \\ $N_\mathrm{ion}$ & $10^{54.54}\,\mathrm{s}^{-1}$ \enddata \label{table:andreas} \end{deluxetable} \subsection{The star-formation and core-collapse supernova rates in NGC\,1614}\label{sec:sedfit} The main goal of this section is to determine two of the most important parameters of any starburst, which are its star-formation rate (SFR) and its core-collapse supernova rate. The (constant) CCSN rate, \ccsnrate, can be related to the (constant) $\mathrm{SFR}$ as follows \citep{perez-torres09a}: \begin{equation} \ccsnrate = \mathrm{SFR}\, \left( \frac{\alpha-2}{\alpha-1} \right) \left( \frac{m_{\rm SN}^{1-\alpha} - m_u^{1-\alpha}}{m_l^{2-\alpha} - m_u^{2-\alpha}} \right) \label{eq:ccsnrate} \end{equation} where $\mathrm{SFR}$ is the (constant) star formation rate in \msunyr, $m_l$ and $m_u$ are the lower and upper mass limits of the initial mass function (IMF, $\Phi \propto m^{-\alpha}$), and $m_{\rm SN}$ is the minimum mass of stars that yield supernovae, assumed to be 8 \msun \citep[e.g.,][]{smartt09}. \citet{mattila01} found an empirical relationship between \lfir and \ccsnrate: \ccsnrate \ $\approx 2.7 \times 10^{-12}$\,(\lir/\Lsun)\,yr$^{-1}$. This implies a CCSN rate for the circumnuclear starburst of NGC\,1614 of $\approx1.08$\,SN\,yr$^{-1}$, for L$_{\rm IR} \simeq 4.0\times 10^{11} L_\odot$, which according to Eq. \ref{eq:ccsnrate} corresponds to a (constant) SFR of $\approx52.9\msunyr$. However, a constant star-formation process is likely to be, for LIRGs in general, and for NGC\,1614 in particular, a poor approximation to the actual starburst scenario \citep{alonso-herrero01}. The opposite case to a constant SFR is that of a single instantaneous starburst. For example, \citet{rosenberg12} used Starburst~99 \citep{leitherer99} to model the emission of NGC\,1614 within an instantaneous starburst scenario, and obtained an average age for the starburst of 6.4\,Myr and an integrated SN rate of 0.9\,yr$^{-1}$. \citet{u12} found a star formation rate of $\mathrm{SFR}_\mathrm{UV+IR} \simeq 51.3\,M_\odot\,\mathrm{yr}^{-1}$. \citet{alonso-herrero01} modeled the star formation of NGC\,1614 using two Gaussian bursts, each of them with a FWHM of 5\,Myr, separated by 5\,Myr, obtaining an age of $\sim11$\,Myr after the peak of the first burst, i.e., a total age of $\sim16$\,Myr. From the extinction corrected [Fe\,{\sc ii}], they predicted a supernova rate of 0.3\,SN/yr$^{-1}$. An intermediate approach is that of an exponentially decaying starburst, which is the approach we have followed here. Namely, we modeled the near-IR to sub-millimeter spectral energy distribution (SED) of NGC\,1614 by combining pure starburst models from \citet{efstathiou00}, revised by \citet{efstathiou09}, and models for the host galaxy. The latter models the emission from the stars using the models by \citet{bruzual03} and the emission from diffuse (cirrus) dust using the model of \citet{efstathiou09}. The fit is shown in Fig.~\ref{fig:andreas}, where photometric data points were obtained from NED\footnote{The NASA/IPAC Extragalactic Database (NED) is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration.}, \citet{soifer01}, \citet{skrutskie06}, and \citet{ishihara10}. The best fit starburst model yields a bolometric luminosity of $10^{11.39}L_\odot$, an initial star formation rate of 85.1 \msunyr (57.7 \msunyr averaged over the duration of the starburst), a core-collapse supernova rate of 0.43\,SN\,yr$^{-1}$, and an ionizing photon flux of $3.47\times10^{54}$\,s$^{-1}$. \subsection{Is there an AGN in the center of NGC\,1614?}\label{sec:agn} The nuclear region, N, shows a non-thermal spectral index of $\alpha \sim -1.80$. Such a steep spectral index seems to be at odds with an AGN origin for the radio emission of region N, but can be reconciled with the scenario of a compact starburst, powered by supernovae and supernova remnants. While such a high value of $\alpha$ is not rare or extreme for starbursts \citep[see, e.g., NGC\,253 in][]{heesen11}, it implies heavy synchrotron losses and, given the large radiation field in the nuclear region (see the prominent 8.7\,$\mu$m continuum and Pa$\alpha$ line emission), also large inverse Compton losses. We note that the radio spectral index is the average value over region N, of $\sim90$\,pc in radius, so in principle we cannot rule out completely the existence of a hidden AGN inside that region, as found in other LIRGs, e.g., in Arp~299-A \citep{perez-torres10}. However, even if there is an AGN, its radio luminosity would not contribute more than $\sim6\%$ and $\sim 7.6\%$ at 3.6 and 6.0\,cm, respectively (see Table~\ref{table:variability}). For comparison, the AGN in Arp~299-A, as found from VLBI observations has a 5.0\,GHz flux density at cm-wavelenghts of about 820\,$\mu$Jy/b \citep{perez-torres10}, which corresponds to a luminosity of $\nu\,L_{\nu, \rm AGN} \sim 9\times 10^{36}$\ergs, and accounts for no more than about 11\% of the compact VLBI 5.0 GHz flux. This luminosity value is also less than 1\% of the total 5.0 GHz emission, as traced by eMERLIN within the region where all SNe/SNRs are exploding \citep[see Fig.~4 and sect.~3.3 in][]{bondi12}. The corresponding 5.0 GHz luminosity of a putative AGN in NGC\,1614 is therefore no more than $\nu\,L_{\nu, \rm AGN} \sim 7.1\times 10^{37}$\ergs\ and, if most of this luminosity comes in turn from a nuclear starburst, the AGN must be even fainter. The ratio of the 8.7\,$\mu$m/Pa$\alpha$ emission also suggests that region N is powered by a burst of star formation. In fact, the ratio in the central $\sim$90 pc agrees well with the ratios obtained for nuclei of H\,{\sc ii} systems, and is significantly lower than obtained for the nuclei of Sy/Sy 2 systems (Table 3 in D\'iaz-Santos et al. 2008), which are known to host an AGN. Similarly, the \emph{Chandra} X-ray emission supports a starburst driven scenario for the central regions of NGC\,1614. We calculated hardness ratios, which are model-independent, for different apertures. We defined the hardness ratio as $\mathrm{HR}=(H-S)(H+S)$, being $H$ and $S$ the hard [2.0--10.0]\,keV and soft [0.5--2.0]\,keV bands, respectively. For an aperture of 3\asec in radius, we get HR$=-0.40$, while for an aperture of 0.3\asec (coincident with region N) we get a significantly harder spectrum, HR=+0.69. The point at which HR becomes positive (i.e., the hard emission dominates) is $\lsim0.4\asec$. The weak hard X-ray emission could easily be due to the presence of X-ray binaries in a compact starburst in the central 0.3\asec (i.e., $\sim110$\,pc), in agreement with the sizes of the starbursts seen in, e.g., Arp~299-A \citep{perez-torres09a, bondi12} and Arp~220 \citep{parra07}. While we cannot rule out completely that an AGN makes also some contribution, the presence of emission lines of Mg\,{\sc xi}, Si\,{\sc xiii} and S\,{\sc xv} in the spectrum (see Fig.~\ref{fig:chandraall}) suggests the existence of SNe and/or SNRs. Although the temperature is somewhat higher than expected for SNRs, with typical values of kT=0.5\,keV \citep[see, e.g.,][]{soria03}, they are in good agreement with values obtained for young supernovae \citep[e.g., SN2001gd in NGC\,5033 had kT=1.1 keV;][]{perez-torres05}. Additionally, using the multi-wavelength optical, radio, and soft X-rays diagnostic diagram from \citet{perez-olea96}, region N would be in the starburst dominated region, as can be seen in Fig.~\ref{fig:diagnostic}, with $\log(L_X/L_\mathrm{5GHz}) < 4.77$ and $\log(L_X/L_\mathrm{H\alpha}) < -0.43$. (The inequality accounts for the fact that the radio and H$\alpha$ luminosity are for region N, while the X-ray luminosity corresponds to a 3\asec aperture.) \begin{figure} %\epsscale{.80} \includegraphics[width=\columnwidth]{ngc1614_diagnostic.eps} \caption{Multi-wavelength diagnostic plot discriminating starbursts from AGN. NGC\,1614 is plotted as a blue star. Since the used X-ray aperture (3\asec) is larger than the H$\alpha$ and 5\,GHz ones (region N, 0.6\asec), NGC\,1614 real position in the diagram will necessarily move down and to the left, making NGC\,1614 fall clearly in the starburst dominated region. We derive an X-ray luminosity in the range [0.1--2.4]\,keV of $10^{40.90}$\,erg/s. Adapted from \citet{perez-olea96}} \label{fig:diagnostic} \end{figure} We also used archival data from \emph{Spitzer} IRS to check the high-resolution spectra for NGC\,1614 looking for [Ne\,{\sc v}] lines at $14.3$ and $24.3\,\mu$m, which would be indicative of the existence of an AGN \citep{genzel98, armus07}, but found no evidence of their presence. Finally, we also fitted the multi-wavelength SED with a combination of starburst models and AGN torus models \citep{efstathiou95} and found that the contribution of the AGN to the total bolometric luminosity, if any, would be at most $\sim10\%$. In summary, all evidence shows that the bulk of the observed emission (at all wavelengths) from the circumnuclear region of NGC\,1614 can be explained with the existence of a powerful starburst, without any need to advocate the existence of an AGN.
14
3
1403.7201
The Luminous Infrared Galaxy NGC 1614 hosts a prominent circumnuclear ring of star formation. However, the nature of the dominant emitting mechanism in its central ~100 pc is still under debate. We present sub-arcsecond angular resolution radio, mid-infrared, Paα, optical, and X-ray observations of NGC 1614, aimed at studying in detail both the circumnuclear ring and the nuclear region. The 8.4 GHz continuum emission traced by the Very Large Array and the Gemini/T-ReCS 8.7 μm emission, as well as the Paα line emission, show remarkable morphological similarities within the star-forming ring, suggesting that the underlying emission mechanisms are tightly related. We used a Hubble Space Telescope/NICMOS Paα map of similar resolution to our radio maps to disentangle the thermal free-free and non-thermal synchrotron radio emission, from which we obtained the intrinsic synchrotron power law for each individual region within the central kiloparsec of NGC 1614. The radio ring surrounds a relatively faint, steep-spectrum source at the very center of the galaxy, suggesting that the central source is not powered by an active galactic nucleus (AGN), but rather by a compact (r &lt;~ 90 pc) starburst (SB). Chandra X-ray data also show that the central kiloparsec region is dominated by SB activity, without requiring the existence of an AGN. We also used publicly available infrared data to model-fit the spectral energy distribution of both the SB ring and a putative AGN in NGC 1614. In summary, we conclude that there is no need to invoke an AGN to explain the observed bolometric properties of the galaxy.
false
[ "SB activity", "SB", "AGN", "NGC", "star formation", "sub-arcsecond angular resolution radio", "a prominent circumnuclear ring", "its central ~100 pc", "the central kiloparsec region", "X-ray observations", "debate", "remarkable morphological similarities", "Paα", "similar resolution", "The radio ring", "the central source", "detail", "the underlying emission mechanisms", "the Paα line emission", "Chandra X-ray data" ]
12.569157
8.296754
-1
482997
[ "Regály, Zs.", "Király, S.", "Kiss, L. L." ]
2014ApJ...785L..31R
[ "Asymmetric Fundamental Band CO Lines as a Sign of an Embedded Giant Planet" ]
11
[ "Konkoly Observatory, Research Center for Astronomy and Earth Sciences, P.O. Box 67, H-1525 Budapest, Hungary; ELTE Gothard-Lendület Research Group, H-9704 Szombathely, Szent Imre Herceg u. 112, Hungary;", "Konkoly Observatory, Research Center for Astronomy and Earth Sciences, P.O. Box 67, H-1525 Budapest, Hungary", "Konkoly Observatory, Research Center for Astronomy and Earth Sciences, P.O. Box 67, H-1525 Budapest, Hungary; Sydney Institute for Astronomy, School of Physics A28, University of Sydney, NSW 2006, Australia" ]
[ "2015A&A...573A...9P", "2015ApJ...799...43H", "2015ApJ...810...94L", "2015RAA....15.1945S", "2016MNRAS.458.3739B", "2017A&A...598A.118C", "2017A&A...601A..24R", "2017ApJ...851...89R", "2017MNRAS.466.4053J", "2018cwla.conf...49M", "2022AJ....164..136S" ]
[ "astronomy" ]
3
[ "accretion", "accretion disks", "hydrodynamics", "line: formation", "methods: numerical", "protoplanetary disks", "techniques: spectroscopic", "Astrophysics - Solar and Stellar Astrophysics", "Astrophysics - Earth and Planetary Astrophysics" ]
[ "1973A&A....24..337S", "1977Ap&SS..51..153W", "1981PThPS..70...35H", "1984ApJ...285...89D", "1986MNRAS.218..761H", "1991ApJ...381..259L", "1994ApJS...95..535G", "1995A&A...297..273H", "1996Icar..124...62P", "1997ApJ...490..368C", "1997Sci...276.1836B", "2000A&A...358..593S", "2000A&AS..141..165M", "2001AJ....121.2065H", "2004ApJ...606L..73B", "2004SPIE.5492.1218K", "2005A&A...434..343A", "2005Icar..179..415H", "2005MNRAS.359..663D", "2006A&A...447..369K", "2006MNRAS.370..529D", "2007ApJ...662.1067H", "2008A&A...486..617K", "2008A&A...487..671K", "2008ApJ...684.1323P", "2008MNRAS.386..973P", "2009A&A...504..461F", "2009A&A...508.1493M", "2009ApJ...699..330S", "2009Icar..199..338L", "2010A&A...523A..69R", "2010Icar..209..616M", "2011A&A...528A..93R", "2012A&A...539A..18M", "2012A&A...539A..98M", "2012A&A...546A..99K", "2012ApJ...749..153J", "2012MNRAS.419.1701R", "2013A&A...556A.148P", "2013ApJ...770...94B", "2013MNRAS.433.2626R" ]
[ "10.1088/2041-8205/785/2/L31", "10.48550/arXiv.1403.2539" ]
1403
1403.2539_arXiv.txt
Protoplanetary disks are expected to emit symmetric double-peaked molecular lines in infrared wavelengths as a result of the Keplerian angular velocity distribution of gas parcels \citep{HorneMarsh1986}. Contrary to this, asymmetric line profiles have been observed in the fundamental band of CO for several young ($1\textrm{--}5$\,Myr) protoplanetary disks \citep{Pontoppidanetal2008,BlakeBoogert2004,Dentetal2005,Salyketal2009,Brownetal2013}. According to the grid-based numerical simulations of \cite{KleyDirksen2006}, disk eccentricity can be excited locally, near the gap opened by an embedded giant planet. As the orbital velocity of the gas parcels is highly supersonic in accretion disks, gas parcels in eccentric orbits result in supersonic velocity deviations in comparison to the circular Keplerian case. \cite{Horne1995} has shown that supersonic turbulence can cause observable line profile distortions in the molecular spectra of protoplanetary disks. \citet{Regalyetal2010} have demonstrated that the double-peaked Keplerian line profiles in the fundamental band of CO are distorted due to the excitation of disk eccentricity in the vicinity of the gap carved by a close orbiting ($\leq1\,\rm AU$) giant planet ($M_\mathrm{p}>1\,M_{Jup}$), which allows an indirect detection of the planet. The theory of resonant excitation mechanisms in accretion disks of \cite{Lubow1991} predicts that the circumstellar disks of close-separation young binaries become fully eccentric due to the orbiting companion. Locally isothermal numerical simulations have confirmed this \citep{Kleyetal2008,KleyNelson2008,Paardekooperetal2008,Marzarietal2009,Regalyetal2011}. However, only fast cooling and low-mass disks favor the excitation of disk eccentricity with radiative and self-gravitating disk approximations \citep{Marzarietal2012,MullerKley2012}. Radiative three-dimensional SPH simulations of \citet{PicognaMarzari2013} have also revealed $\sim0.1$ eccentricity excitation of circumstellar disks in binaries. Nonetheless, a disk having a non-constant radial eccentricity profile with $\sim0.2\textrm{--}0.3$ average eccentricity emits clearly asymmetric fundamental band CO lines \citep{Regalyetal2011}. Since no close stellar mass companions were found for the majority of disks that emits asymmetric lines, it is worth investigating whether a giant planet is able to excite global disk eccentricity inside its orbit where the CO is thermally excited. We present the excitation of global disk eccentricity by means of two-dimensional hydrodynamical simulations, which leads to the formation of asymmetric double-peaked line profiles of CO.
In this Letter, we presented fundamental band spectra of the CO molecule in circumstellar disks gravitationally perturbed by an embedded $2.5,\,5$, and $10\, M_\mathrm{Jup}$ giant planet orbiting $M_*=0.5,\,1$, and $2\,M_\odot$ stars, respectively, at $1\,\mathrm{AU}\leq a_\mathrm{p}\leq 10\,\mathrm{AU}$. We found that the CO lines are asymmetric due to the development of disk eccentricity inside the region where CO is thermally excited. The principal results of our study are: \begin{enumerate} \itemsep2pt \parskip0pt \parsep0pt \parindent0pt \itemindent0pt \item{A quasi-static, globally eccentric disk state develops inside the planetary orbit with $\langle {e}_\mathrm{disk}\rangle=0.2\textrm{--}0.25$ after several thousand orbits by $t\simeq40\times10^3$\,yr.} \item{As a result, asymmetric lines are formed with magnitude of $\sim10\%\textrm{--}20\%$ measured in red versus blue peak, whose strength increases with decreasing orbital distance of the planet.} \item{The lines are off-centered, i.e., shifted toward the higher flux peak by $\sim4\textrm{--}10\,\mathrm{km\,s^{-1}}$, such that the smaller the orbital distance of the planet or the larger the stellar mass, the greater the shift.} \item{The line asymmetry--$J$ slope informs us whether the giant planet orbits outside (asymmetry decreases with $J$) or inside (asymmetry increases with $J$) the thermally excited CO region. The critical orbital distance is $a_{p}\simeq3,\,5$, and $7\,\rm AU$ for an $M_*=0.5,\,1$, and $2\,M_\odot$ star, where the asymmetry--$J$ slope changes its sign.} \item{For close planetary orbits ($a_{p}\leq1,\,3$, and $5\,\rm AU$ for an $M_*=0.5,\,1$, and $2\,M_\odot$ star), the lines are highly distorted.), the CO lines do not show clear asymmetry; rather, highly perturbed lines are formed (as was shown previously in \citealp{Regalyetal2010}) because the planet-induced gap is formed inside the CO excitation region.} \item{The peak distances, peak positions, and magnitude of the central dip are found to be proportional to $J$ independent of the planetary and stellar masses.} \end{enumerate} In light of our findings (line shape dependency on $J$), we have to emphasize that care must be taken when applying line profile averaging in observational studies: averaging of only neighboring transitions (close in $J$) are suggested in which case the difference in the line shapes are modest. Since the development of disk eccentricity presumably depends on radiative processes as well as the disks physical parameters (e.g., viscosity, aspect ratio, etc.) and naturally the mass of the perturbing planet, further investigation is inevitable before attempting any detailed characterization of planets from the observed line profiles. A young ($1\textrm{--}5$\,Myr), gas-rich disk that emits asymmetric CO lines may harbor giant planets which should be formed early in the disks life. The classical core accretion scenario predicts that the formation of giant planet takes $5\textrm{--}10$\,Myr \citep{Pollacketal1996}, which is very close to or even larger than the observed disk lifetimes \citep{Haischetal2001,Hernandezetal2007}. Recently, several theoretical studies of core accretion have reported a shorter formation time ($1\textrm{--}5$\,Myr) of giant planets. Planetary migration prevents the severe depletion of the feeding zone as observed in in situ calculations; however, Type\,I migration must be artificially slowed by an order of magnitude \citep{Alibertetal2005}. Assuming a protoplanetary disk with a high solid surface density of $10\,\rm g\,cm^{-2}$ and dust opacity in the protoplanet's envelope equals to 2\% that of interstellar material will result in a $2.5\textrm{--}3$\,Myr formation time of a Jupiter-mass planet \citep{Hubickyjetal2005,Lissaueretal2009,Movshovitzetal2010}. However, the required solid surface density is about an order of magnitude larger than that proposed by the minimum solar mass nebula model at 5.2\,AU assuming a canonical $10^{-2}$ dust-to-gas mass ratio \citep{Hayashi1981}. Note that giant planets can be formed much faster within the context of gravitational instability theory, but it requires too high disk masses ($>0.1\,M_\odot$) and an efficient cooling mechanism to operate \citep{Boss1997}. Another possible scenario could be vortex-aided planet formation in which the large-scale anticyclonic vortex developed near the disks outer dead zone edge (\citealp{Regalyetal2012}) may induce planet formation. The core of a giant planet is subject to be trapped temporarily in the vortex vicinity \citep{Regalyetal2013}, where large amounts of dusty material are accumulated. Therefore, gas giants can be formed rapidly in both scenarios. In the core-accretion scenario, the slow growth regime of the giant planetary core might be circumvented by shortening the slow growth phase within the isolation ($\sim10\,M_\oplus$) and the runaway growth ($\sim20\,M_\oplus$) mass limits. In the gravitational instability scenario, the disk might be gravitationally unstable without requiring high disk mass due to the large amount of dusty and gaseous material accumulated in the planetary trap. As the presented spectroscopic phenomenon can be applied for the characterization of planets still embedded in their host disks, it can deliver new types of empirical tests of planet formation theories.
14
3
1403.2539
We investigate the formation of double-peaked asymmetric line profiles of CO in the fundamental band spectra emitted by young (1-5 Myr) protoplanetary disks hosted by a 0.5-2 M <SUB>⊙</SUB> star. Distortions of the line profiles can be caused by the gravitational perturbation of an embedded giant planet with q = 4.7 × 10<SUP>-3</SUP> stellar-to-planet mass ratio. Locally isothermal, two-dimensional hydrodynamic simulations show that the disk becomes globally eccentric inside the planetary orbit with stationary ~0.2-0.25 average eccentricity after ~2000 orbital periods. For orbital distances 1-10 AU, the disk eccentricity is peaked inside the region where the fundamental band of CO is thermally excited. Hence, these lines become sensitive indicators of the embedded planet via their asymmetries (both in flux and wavelength). We find that the line shape distortions (e.g., distance, central dip, asymmetry, and positions of peaks) of a given transition depend on the excitation energy (i.e., on the rotational quantum number J). The magnitude of line asymmetry is increasing/decreasing with J if the planet orbits inside/outside the CO excitation zone (R <SUB>CO</SUB> &lt;= 3, 5, and 7 AU for a 0.5, 1, and 2 M <SUB>⊙</SUB> star, respectively), thus one can constrain the orbital distance of a giant planet by determining the slope of the peak asymmetry-J profile. We conclude that the presented spectroscopic phenomenon can be used to test the predictions of planet formation theories by pushing the age limits for detecting the youngest planetary systems.
false
[ "planet formation theories", "planet", "orbital distances", "orbital periods", "CO", "line asymmetry", "J", "an embedded giant planet", "a giant planet", "peaks", "the CO excitation zone", "the embedded planet", "wavelength", "double-peaked asymmetric line profiles", "the youngest planetary systems", "the orbital distance", "the fundamental band spectra", "the disk eccentricity", "the peak asymmetry-J profile", "the planet" ]
9.51346
13.079536
-1
526821
[ "Errmann, R.", "Raetz, St.", "Kitze, M.", "Neuhäuser, R.", "the YETI team" ]
2014arXiv1403.6031E
[ "The search for transiting planets using the YETI network" ]
1
[ "-", "-", "-", "-", "-" ]
[ "2019A&A...624A.110F" ]
[ "astronomy" ]
4
[ "Astrophysics - Earth and Planetary Astrophysics", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "2002AJ....124.1585C", "2005AJ....129..907B", "2005AN....326..134B", "2007ApJ...661.1119B", "2008hsf1.book..136K", "2011AN....332..547N", "2012ApJ...755...42V", "2013AN....334..673E", "2013ApJ...774...53B" ]
[ "10.48550/arXiv.1403.6031" ]
1403
1403.6031_arXiv.txt
\label{intr} Young transiting planets could play a key role to distinguish planet formation scenarios, as it is possible to test evolutionary models with the well determined parameters of such planets, like radius, mass, and age. As no transiting planets with ages younger than $100$\,Myr were found so far, we monitor young open clusters to search for transit signals among young stars. To increase phase coverage, continuous observations are needed. Therefore we established YETI (Young Exoplanet Transit Initiative), a network of ground-based telescopes with mirror diameters of $0.4$ to $2$\,m (see Neuh{\"a}user {\it et al.} \cite{neu11} for further details). Fig.\,\ref{Fig:YETI_map} shows a map with all telescope sites as well as the telescope sizes. Each cluster is observed for at least three years with three campaign runs of lengths longer than one week. Observations are done in the $R$-band. \begin{figure}[t] \centerline{\includegraphics[width=0.98\textwidth,clip=]{YETI_map1.eps}} \caption{Map of the YETI telescopes. With the network it is possible to observe continuously.} \label{Fig:YETI_map} \vspace*{-5.mm} \end{figure} The first two target clusters of YETI were Trumpler\,37 (Tr37) and 25\,Ori. Some properties of both clusters are summarized in Tab.\,\ref{Tab:cluster_prop}. The monitoring of Tr37 started in summer 2009 at the University Observatory Jena. Data of the YETI telescopes were gathered in summers 2010 and 2011, each year three campaign runs were performed. The observations of 25\,Ori at the University Observatory Jena started in January 2010. In this first phase of the observation two additional telescopes (Gunma/Japan and CIDA/Venezuela) joined the photometric monitoring in January and February 2010. 25\,Ori became a target of the YETI project in the winter seasons 2010/2011, 2011/2012, and 2012/2013. Fifteen different observatories, spread worldwide at different longitudes, participated in the campaigns. Data reduction (dark/bias and flat-field correction) as well as aperture photometry with adjusted aperture to the seeing conditions was done night by night and telescope by telescope. For a particular star, the final lightcurve was created by doing differential photometry (Broeg {\it et al.} \cite{bro05}) for all nights and telescopes on a small sample of comparison stars (similar brightness and color, small angular distance). A detailed description can be found in Errmann {\it et al.} (\cite{err13b}). \begin{table} \small \begin{center} \caption{Properties of the stars in the clusters Trumpler\,37 and 25\,Ori.} \label{Tab:cluster_prop} \vspace*{-3.mm} \begin{tabular}{lcc} \hline\hline & Trumpler\,37 & 25\,Ori \\ \hline Age [Myr] & $\sim4$ [1] & 7-10 [2] \\ Distance [pc] & $\sim870$ [3] & $\sim330$ [4] \\ Number member stars & 774 [5] & $\sim250$ [6] \\ \hline\hline \multicolumn{3}{c}{ [1]~Kun {\it et al.} (\cite{kkb08}); [2]~Brice{\~n}o {\it et al.} (\cite{bri07}); [3]~Contreras {\it et al.} (\cite{con02}); }\\ \multicolumn{3}{c}{ [4]~Brice{\~n}o {\it et al.} (\cite{bri05}); [5]~Errmann {\it et al.} (\cite{err13a}); [6]~Brice{\~n}o {\it et al.} (\cite{bri13}) } \end{tabular} \end{center} \vspace*{-5.mm} \end{table} Fig.\,\ref{Fig:timecoverage} shows the time coverage of the observations for a star in Tr37 for the third YETI campaign in 2011. As the star is located in a larger separation to the cluster center, it does not fit into the field of view of some telescopes. Furthermore, not all telescopes could always allocate time and some suffer from seasonal bad weather, hence only 5 of the YETI sites are present. But even with that smaller number of telescopes, we could observe nearly continuously for $48.5$\,h at the end of the campaign. Fig.\,\ref{Fig:phasecoverage} shows the phase coverage of the same star. % For a period up to 10\,d we reach 100\% phase coverage and even for a period of 50\,d the phase coverage is better than 70\%. The coverage at periods of a multiple of a full day is slightly worse, as a telescope in the pacific ocean is missing. \begin{figure} \centerline{\includegraphics[height=0.93\textwidth,angle=270]{3218-2a-60sergR31-1-colorcoded_sorted_sub-coverage_sw.eps}} \caption{Observation times of the first transiting candidate in Trumpler\,37 during the third YETI campaign in 2011.} \label{Fig:timecoverage} \vspace*{-5.mm} \end{figure} \begin{figure} \centerline{\includegraphics[height=0.93\textwidth,angle=270]{phase_coverage_3218.eps}} \caption{Phase coverage of the first transiting candidate in Trumpler\,37.} \label{Fig:phasecoverage} \vspace*{-5.mm} \end{figure}
14
3
1403.6031
To search for young transiting planets in continuous light curves, we monitor young open clusters (2-200 Myr) with the YETI network. Here we report the first transiting candidates (two in Trumpler 37, one in 25 Ori). Follow-up observations of the candidates are partly done.
false
[ "young open clusters", "continuous light curves", "young transiting planets", "YETI", "the YETI network", "Trumpler", "the first transiting candidates", "first", "25 Ori", "the candidates", "Trumpler 37", "Follow-up observations", "2-200 Myr", "2", "25", "one", "two", "we" ]
6.71333
13.59005
-1
437669
[ "Guennou, L.", "Biviano, A.", "Adami, C.", "Limousin, M.", "Lima Neto, G. B.", "Mamon, G. A.", "Ulmer, M. P.", "Gavazzi, R.", "Cypriano, E. S.", "Durret, F.", "Clowe, D.", "LeBrun, V.", "Allam, S.", "Basa, S.", "Benoist, C.", "Cappi, A.", "Halliday, C.", "Ilbert, O.", "Johnston, D.", "Jullo, E.", "Just, D.", "Kubo, J. M.", "Márquez, I.", "Marshall, P.", "Martinet, N.", "Maurogordato, S.", "Mazure, A.", "Murphy, K. J.", "Plana, H.", "Rostagni, F.", "Russeil, D.", "Schirmer, M.", "Schrabback, T.", "Slezak, E.", "Tucker, D.", "Zaritsky, D.", "Ziegler, B." ]
2014A&A...566A.149G
[ "Mass profile and dynamical status of the z ~ 0.8 galaxy cluster LCDCS 0504" ]
13
[ "Astrophysics and Cosmology Research Unit, University of KwaZulu-Natal, Durban, 4041, South Africa ; Aix-Marseille Université, CNRS, LAM (Laboratoire d'Astrophysique de Marseille) UMR 7326, 13388, Marseille, France", "INAF/Osservatorio Astronomico di Trieste, via Tiepolo 11, 34143, Trieste, Italy; Institut d'Astrophysique de Paris (UMR 7095: CNRS &amp; UPMC), 98bis Bd Arago, 75014, Paris, 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", "Departamento de Astronomia, Instituto de Astronomia Geofìsica e Ciências Atmosfèricas, Universidade de São Paulo, Rua do Matão 1226, 05508-900, São Paulo, Brazil", "Institut d'Astrophysique de Paris (UMR 7095: CNRS &amp; UPMC), 98bis Bd Arago, 75014, Paris, France", "Department Physics &amp; Astronomy &amp; Center for Interdisciplinary Exploration and Research in Astrophysics (CIERA), Northwestern University, Evanston, IL, 60208-2900, USA", "Institut d'Astrophysique de Paris (UMR 7095: CNRS &amp; UPMC), 98bis Bd Arago, 75014, Paris, France", "Departamento de Astronomia, Instituto de Astronomia Geofìsica e Ciências Atmosfèricas, Universidade de São Paulo, Rua do Matão 1226, 05508-900, São Paulo, Brazil", "Institut d'Astrophysique de Paris (UMR 7095: CNRS &amp; UPMC), 98bis Bd Arago, 75014, Paris, France", "Department of Physics and Astronomy, Ohio University, 251B Clippinger Lab, Athens, OH, 45701, USA", "Aix-Marseille Université, CNRS, LAM (Laboratoire d'Astrophysique de Marseille) UMR 7326, 13388, Marseille, France", "Fermi National Accelerator Laboratory, PO Box 500, Batavia, IL, 60510, USA; CSC/STSCi, 3700 San Martin Dr., Baltimore, MD, 21218, USA", "Aix-Marseille Université, CNRS, LAM (Laboratoire d'Astrophysique de Marseille) UMR 7326, 13388, Marseille, France", "OCA, Cassiopée, Boulevard de l'Observatoire, BP 4229, 06304, Nice Cedex 4, France", "INAF/Osservatorio Astronomico di Bologna, via Ranzani 1, 40127, Bologna, Italy; OCA, Cassiopée, Boulevard de l'Observatoire, BP 4229, 06304, Nice Cedex 4, France", "23 rue d'Yerres, 91230, Montgeron, France", "Aix-Marseille Université, CNRS, LAM (Laboratoire d'Astrophysique de Marseille) UMR 7326, 13388, Marseille, France", "Fermi National Accelerator Laboratory, PO Box 500, Batavia, IL, 60510, USA", "Aix-Marseille Université, CNRS, LAM (Laboratoire d'Astrophysique de Marseille) UMR 7326, 13388, Marseille, France", "Steward Observatory, University of Arizona, 933, N. Cherry Ave. Tucson, AZ, 85721, USA; Department of Astronomy &amp; Astrophysics, University of Toronto, 50 St George Street Toronto, Ontario, M5S 3H4, Canada", "Fermi National Accelerator Laboratory, PO Box 500, Batavia, IL, 60510, USA", "Instituto de Astrofísica de Andalucía (CSIC), Glorieta de la Astronomía s/n, 18008, Granada, Spain", "Kavli Institute for Particle Astrophysics and Cosmology, Stanford University, 2575 Sand Hill Rd., Menlo Park, CA, 94025, USA", "Institut d'Astrophysique de Paris (UMR 7095: CNRS &amp; UPMC), 98bis Bd Arago, 75014, Paris, France", "OCA, Cassiopée, Boulevard de l'Observatoire, BP 4229, 06304, Nice Cedex 4, France", "Aix-Marseille Université, CNRS, LAM (Laboratoire d'Astrophysique de Marseille) UMR 7326, 13388, Marseille, France", "Department of Physics and Astronomy, Ohio University, 251B Clippinger Lab, Athens, OH, 45701, USA", "Laboratório de Astrofísica Teórica e Observacional, Universidade Estadual de Santa Cruz, Ilhéus, Brazil", "OCA, Cassiopée, Boulevard de l'Observatoire, BP 4229, 06304, Nice Cedex 4, France", "Aix-Marseille Université, CNRS, LAM (Laboratoire d'Astrophysique de Marseille) UMR 7326, 13388, Marseille, France", "Gemini Observatory, Casilla 603, La Serena, Chile; Argelander-Institut für Astronomie, Universitët Bonn, Auf dem Hügel 71, 53121, Bonn, Germany", "Physics Department, University of California, Santa Barbara, CA, 93601, USA", "OCA, Cassiopée, Boulevard de l'Observatoire, BP 4229, 06304, Nice Cedex 4, France", "Fermi National Accelerator Laboratory, PO Box 500, Batavia, IL, 60510, USA", "Steward Observatory, University of Arizona, 933, N. Cherry Ave. Tucson, AZ, 85721, USA", "University of Vienna, Department of Astrophysics, Türkenschanzstr. 17, 1180, Wien, Austria" ]
[ "2015MNRAS.452..944B", "2016A&A...594A..51B", "2016A&A...595A..30V", "2017A&A...602A..20B", "2017A&A...607A..81B", "2017AJ....154..113T", "2017MNRAS.468..662R", "2018arXiv181010474T", "2019SSRv..215...25P", "2020ApJ...888..106T", "2021A&A...650A.105B", "2024ApJ...965..117B", "2024arXiv240208453K" ]
[ "astronomy" ]
4
[ "galaxies: clusters: general", "galaxies: kinematics and dynamics", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
[ "1976A&A....49..137C", "1982AJ.....87..945K", "1986desd.book.....S", "1987A&A...183..217R", "1987gady.book.....B", "1988AJ.....95..985D", "1989GeCoA..53..197A", "1990AJ....100...32B", "1991A&A...246L..10B", "1992ApJ...400..699B", "1993A&AS...97..443K", "1993ApJ...413..492M", "1994ApJ...423..539E", "1995ApJ...438L.115L", "1996A&A...309...65S", "1996A&A...311...95B", "1996A&A...313..697B", "1996MNRAS.279..349D", "1997ApJ...481..633D", "1997ApJ...490..493N", "1998ApJ...503..518F", "1999MNRAS.309..610D", "2000AJ....119.2038V", "2001A&A...365L..74N", "2001ApJ...563..629N", "2001MNRAS.321..155L", "2001Natur.409...39B", "2002A&A...386..816B", "2002ASPC..268..395K", "2002ASSL..272...39G", "2002ApJ...567L..27M", "2003ApJ...585..205B", "2003MNRAS.343..401L", "2004A&A...424..779B", "2004A&A...427..397H", "2004ApJ...600..657K", "2004ApJ...604..604O", "2004ApJ...606..819M", "2004ApJ...613...95C", "2004MNRAS.347.1198S", "2005A&A...433..777P", "2005A&A...434..433B", "2005A&A...440..775K", "2005A&A...443...17A", "2005ApJ...618...46J", "2005ApJ...629..781K", "2005MNRAS.356..309L", "2005MNRAS.363..705M", "2006A&A...452...75B", "2006A&A...456...23B", "2006MNRAS.367.1463L", "2006MNRAS.369.2013R", "2006MNRAS.371.1777J", "2007A&A...461..881L", "2007A&A...466..437W", "2007A&A...470...39R", "2007A&A...476L...1T", "2007ApJ...668..643L", "2007NJPh....9..447J", "2007arXiv0710.5636E", "2008ApJ...681..151C", "2008MNRAS.383..879A", "2009ApJ...700.1603F", "2009MNRAS.399..812W", "2010A&A...514A..32B", "2010A&A...520A..30M", "2010A&A...523A..21G", "2010ApJ...711.1033Z", "2010ApJ...723.1678C", "2010MNRAS.408.2442W", "2011MNRAS.412..800S", "2011MNRAS.413..789L", "2013A&A...551A..23E", "2013A&A...558A...1B", "2013A&A...559A..89W", "2013ApJ...764...58G", "2013ApJ...766...32B", "2013ApJ...773..116G", "2013ApJS..208...19H", "2013MNRAS.428.2885D", "2013MNRAS.428.2921D", "2013MNRAS.429.3079M", "2014A&A...561A.112G", "2014A&A...571A..16P", "2014A&A...571A..20P" ]
[ "10.1051/0004-6361/201322447", "10.48550/arXiv.1403.3614" ]
1403
1403.3614_arXiv.txt
\label{s:intro} The study and characterization of the internal dynamics of galaxy clusters is an important way to understand their evolutionary history, which is itself related to the evolutionary history of the universe. The most classical way to characterize the dynamics of clusters is through the analysis of the projected phase space distribution of their member galaxies, e.g. via methods based on the Jeans equation \citep{BT87}, such as the Dispersion-Kurtosis \citep{LM03}, distribution-function \citep{Wojtak+09} and MAMPOSSt \citep{MBB13} methods, or the Caustic method calibrated on numerical simulations \citep{DG97}. All these methods assume spherical symmetry and most of them (except the Caustic method) also assume dynamical relaxation of the cluster. These methods have been applied to several nearby (and massive) clusters of galaxies \citep[see][]{KG82,vanderMarel+00,BG03,LM03,BK04,KBM04,Biviano06b,Lokas+06,WL10}. \begin{figure*} \begin{center} \includegraphics[width=13cm]{HST_caustics.ps} \end{center} \caption{HST image of the core of LCDCS~0504. The size of the field is 38$\times$34 arcsec$^{2}$, corresponding to $285 \times 255$~kpc$^{2}$ at $z = 0.794$. Multiple imaged systems used in this work are labeled. From the best fit strong lensing model, we draw in red the tangential critical curve at z=3 and the corresponding caustic lines in orange} \label{f:multiple} \end{figure*} Given that clusters formed relatively recently according to the hierarchical scenario of structure evolution in the universe \citep[e.g.][]{BG01}, accretion of matter from the surrounding field, in the form of galaxy groups, complicate their internal structure. Detection of secondary structures, or substructures, in clusters is obtained using other methods, either based on the projected distributions of cluster galaxies \citep[e.g.][]{DS88,Escalera+94,Biviano+96,SG96,Barrena+02,GB02,Ramella+07} or on X-ray data for the intra-cluster gas \citep{Briel+91,MFG93,Neumann+01,OHara+04,PBF05,Bohringer+10}. Detection and characterization of these substructures is a direct way to constrain the cluster building history \citep[e.g.][ and references therein]{Adami+05}. These last years, the characterization of the mass distribution and substructures of galaxy clusters has been made possible by investigating deep and high quality data that enable the measurement of weak lensing signal and the detection of strong gravitationally lensed features \citep[e.g.][]{Cypriano+04,Markevitch+04,Bardeau+05,Jee+05,Coe+10,LKG11}. It is still relatively uncommon to see cluster dynamical studies based simultaneously on the Jeans analysis, and on the X-ray and lensing data, especially for high redshift clusters. This is due to the extreme difficulty in obtaining both deep and high resolution X-ray imaging, deep optical and infrared imaging, and faint galaxy spectroscopy. As a consequence, our information on the internal structure and dynamics of distant clusters is still relatively limited. In this paper, we perform a detailed study of the internal structure and dynamics of the rich cluster LCDCS~0504 at redshft $z=0.7943$, also known as Cl~J1216.8-1201 \citep{Nelson+01}, using simultaneously spectroscopic optical data for cluster galaxies, as well as X-ray and strong lensing (SL, hereafter) data. This cluster is part of the DAFT/FADA survey \citep{Guennou+10} and the analysis presented here is a proof of concept for similar analysis to be performed on other clusters of the DAFT/FADA sample. In Sect.~\ref{s:data} we present our data-set. Our SL determination of the cluster mass distribution is described in Sect.~\ref{s:sl}. In Sect.~\ref{s:x} we use the X-ray emission from the hot intra-cluster medium (ICM) to constrain the cluster mass profile. This is also determined using galaxies as tracers in Sect.~\ref{s:kin}. We compare the different mass profile determinations in Sect.~\ref{s:cmp}. In Sect.~\ref{s:bary} we analyse the cluster hot gas mass fraction. We discuss our results in Sect.~\ref{s:summ}, where we also draw our conclusions. Throughout this paper we adopt $H_0=70$ km~s$^{-1}$~Mpc$^{-1}$, $\Omega_m=0.3$, $\Omega_{\Lambda}=0.7$. In this cosmology, 1 arcmin corresponds to 449 kpc at the cluster redshift. \begin{table} \begin{center} \caption{Available data for the LCDCS~0504 cluster.} \label{t:datasumm} \begin{tabular}{lcc} \hline & Archival data & DAFT/FADA data \\ \hline & & \\ Optical imaging & $VRIz$ (VLT/FORS2) & B (Blanco/MOSAIC) \\ & $F814W$ (HST/ACS) & \\ & & \\ IR imaging & Spitzer/IRAC1 and 2 & \\ & & \\ Optical & VLT/FORS2 & Gemini/GMOS \\ spectroscopy & & \\ & & \\ X-ray imaging & XMM-Newton & \\ & (PN/MOS1/MOS2) & \\ \hline \smallskip \end{tabular} \end{center} \end{table}
\label{s:summ} We have analyzed the mass profile $M(r)$ of a $z \approx 0.8$ cluster with the SL technique, using the X-ray emission from the intra-cluster hot gas, and using galaxies as tracers of the gravitational potential. The different determinations of the cluster $M(r)$ disagree, especially in the inner regions. The SL $M(r)$ is slightly above but still consistent with the kinematic determination, but both are significantly above the X-ray $M(r)$ determination. This discrepancy is unlikely to be caused by an unrelaxed dynamical status of the cluster. This could cause an overestimate of the cluster velocity dispersion and hence the cluster mass estimate from kinematics \citep[see, e.g.,][]{Biviano+06} and an incomplete thermalization of the intra-cluster gas, leading to an underestimate of the cluster mass estimates from X-ray \citep[e.g.][]{Rasia+06}, but it would not affect the lensing mass estimate. Moreover, an unrelaxed dynamical status is not supported by the analyses of substructures by \citet{Guennou+13}. In that paper, we have used the \citet[][SG hereafter]{SG96} hierarchical method for the detection of substructures in the distribution of galaxies and searched for substructures in the X-ray data (described in Sect.~\ref{ss:imaging}), by analysing the residuals of the subtraction of a symmetric elliptical $\beta$-model from the X-ray image (see Guennou et al. 2013 for details). Seven substructures were detected by the SG technique, all with masses below 10\% of the total cluster mass. Of these, only one was also detected in X-rays, with an X-ray luminosity of $\approx 8 \%$ the total cluster X-ray luminosity. This analysis indicates that any major perturbation of the LCDCS~0504 dynamical status must thus have occurred sufficiently long ago for the remnants of the merging groups to have disappeared. Another interesting possibility is that we see the cluster with its major axis along the line-of-sight. This is suggested by the circularly symmetric SL configuration and by the small ellipticity of the cD galaxy, since the elongation of cD galaxies generally reflects those of their host clusters \citep[e.g.][]{RK87,KASL02} (see Fig.~\ref{f:multiple}). It has been shown both on numerical simulations \citep{KE05} and observationally \citep{Wojtak13}, that clusters are prolate not only in position space but also in velocity space, and the major axes of the spatial and velocity distributions are aligned. The orientation of the cluster with the major axis along the line-of-sight then results in an overestimate of the cluster mass estimatd from SL and velocity dispersion. According to \cite{Wojtak13}, the mean ratio of the velocity dispersions along the minor and major axes of a cluster is $\simeq 0.78$. This implies a ratio of the velocity dispersion along the major axis to the mean cluster velocity dispersion of 1.16, i.e. a 32\% mass overestimate at a given radius. This is still not sufficient to remove the systematic difference between the mass profile derived from kinematics and that derived by the X-ray analysis. The alignment effect just discussed could also induce an overestimate of mass profile concentration value. This could explain the disagreement we find with the theoretical predictions of \citet{BHHV13} and \citet{DeBoni+13} (see Fig.~\ref{f:rvrs}). Whatever the cause for the X-ray vs. kinematics and SL $M(r)$ discrepancy, substantial systematic underestimates of cluster masses by the X-ray methodology could be interesting for cosmology, as it could alleviate the tension between the $\sigma_8$ values found by the Planck collaboration using the Cosmic Microwave Background power-spectrum on one hand and cluster counts obtained by the Sunyaev-Zeldovich effect (using X-ray masses as mass calibrators \citealp{Planck13}): If X-ray masses are underestimated at given SZ signal, this means the distribution of SZ counts above a given mass threshold is underestimated, meaning that $\Omega_{\rm m}$ ($\sigma_8$) is underestimated (overestimated), which would bring the best-fit value more in line with the CMB value. Another intriguing result of our analysis is the discovery that the gas mass fraction is anomalously low near the center of the LCDCS~0504 cluster. Given the relaxed, symmetric morphology of the X-ray emission (see Fig.~\ref{fig:LCDCS0504allxmm0.3_7.0}), it is unlikely that this anomaly could be attributed to the effects of a major merger displacing the gas from the center, as in the case of the Bullet cluster \citep{Barrena+02,Markevitch+02}. Alternatively, the gas could have been ejected by AGN outbursts, while the effects of SNe explosions should not be significant \citep{CO08,Dubois+13}. \citet{Dubois+13} predict a 30\% loss in the core due to AGN outflows, which is not to far from our observed deficiency (with respect to the average of other clusters) of \mbox{$\simeq 60$}\% (see Fig.~\ref{f:fgas}), given the large observational uncertainties. The main issue with the AGN hypothesis is that there is no evidence of a radio source in the NVSS catalog or in the X-rays as there is no detectable point source at the location of the cD, although there is a hint of a cool core. Also, we have no evidence of broad lines in the optical spectrum of the cD. All this lack of evidence does, however, tell us, is that the assumed AGN activity have subsided long enough ago so that all strong electromagnetic signatures of AGN activity have now subsided. In the near future, we plan to extend the dynamical and structural analysis presented here to clusters with sufficient spectroscopic information in the full DAFT/FADA cluster set. Expanding our data-sets should allow us to determine if the anomalies identified in LCDCS~0504 are a characteristic of high-$z$ clusters or not. Hopefully, with a larger sample we will be able to unveil the hidden systematics causing discrepant determinations of cluster mass profiles by different methods, and to relate these systematics to the currently not well understood physics of the intra-cluster baryons.
14
3
1403.3614
Context. Constraints on the mass distribution in high-redshift clusters of galaxies are currently not very strong. <BR /> Aims: We aim to constrain the mass profile, M(r), and dynamical status of the z ~ 0.8 LCDCS 0504 cluster of galaxies that is characterized by prominent giant gravitational arcs near its center. <BR /> Methods: Our analysis is based on deep X-ray, optical, and infrared imaging as well as optical spectroscopy, collected with various instruments, which we complemented with archival data. We modeled the mass distribution of the cluster with three different mass density profiles, whose parameters were constrained by the strong lensing features of the inner cluster region, by the X-ray emission from the intracluster medium, and by the kinematics of 71 cluster members. <BR /> Results: We obtain consistent M(r) determinations from three methods based on kinematics (dispersion-kurtosis, caustics, and MAMPOSSt), out to the cluster virial radius, ≃1.3 Mpc and beyond. The mass profile inferred by the strong lensing analysis in the central cluster region is slightly higher than, but still consistent with, the kinematics estimate. On the other hand, the X-ray based M(r) is significantly lower than the kinematics and strong lensing estimates. Theoretical predictions from ΛCDM cosmology for the concentration-mass relation agree with our observational results, when taking into account the uncertainties in the observational and theoretical estimates. There appears to be a central deficit in the intracluster gas mass fraction compared with nearby clusters. <BR /> Conclusions: Despite the relaxed appearance of this cluster, the determinations of its mass profile by different probes show substantial discrepancies, the origin of which remains to be determined. The extension of a dynamical analysis similar to that of other clusters of the DAFT/FADA survey with multiwavelength data of sufficient quality will allow shedding light on the possible systematics that affect the determination of mass profiles of high-z clusters, which is possibly related to our incomplete understanding of intracluster baryon physics. <P />Table 2 is available in electronic form at <A href="http://www.aanda.org/10.1051/0004-6361/201322447/olm">http://www.aanda.org</A>
false
[ "nearby clusters", "other clusters", "strong lensing estimates", "mass profiles", "archival data", "multiwavelength data", "shedding light", "prominent giant gravitational arcs", "≃1.3 Mpc", "sufficient quality", "intracluster baryon physics", "high-z clusters", "the inner cluster region", "the central cluster region", "high-redshift clusters", "Mpc", "deep X", "three different mass density profiles", "kinematics", "various instruments" ]
12.716468
4.714561
-1
597221
[ "Bergemann, Maria", "Kudritzki, Rolf-Peter", "Davies, Ben" ]
2014dapb.book..217B
[ "Spectroscopic Analysis of Cool Giants and Supergiants" ]
2
[ "-", "-", "-" ]
[ "2015A&A...583A.122F", "2015MNRAS.450.1900A" ]
[ "astronomy" ]
6
[ "Astrophysics - Solar and Stellar Astrophysics" ]
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[ "10.1007/978-3-319-06956-2_19", "10.48550/arXiv.1403.3087" ]
1403
1403.3087_arXiv.txt
\label{sec:1} Cool evolved stars are perhaps the most enigmatic cosmic objects with luminosities spanning several orders of magnitude. The stars are at the end stages of stellar evolution (Fig. \ref{fig:1}) occupying the coolest vertical strip on the Herzsprung-Russell diagram, the Hayashi limit for fully convective stars. \begin{figure}[!htb] \sidecaption \includegraphics[scale=.35]{fig1.pdf} \caption{Stellar tracks including the effect of rotation for stellar models with initial mass of $0.8$ to $120$ $M_{\odot}$ \cite[][]{2013A&A...558A.103G}. Cool giants and red supergiants occupy the stripe extending vertically from $1 < \log (L/L_{\odot}) < 6$.} \label{fig:1} \end{figure} Low- and intermediate-mass stars evolve to the red giant branch (RGB) and asymptotic giant branch (AGB) after having spent most of their lifetime on the main sequence. These stars are cool and luminous, with effective temperatures $\teff$ between $2000$ and $5500$ K and luminosities $L$ between $10$ and $10^4$ $L_\odot$. The stars have a wide range of ages, from 1 to $> 10$ Gyr, and thus trace chemical composition of interstellar matter in galaxies now and in the past. High-mass stars are those with masses from $10$ to $60$ $M_\odot$; they evolve and explode quickly. Red supergiants (RSG) are young, typically $< 50$ Myr, yet extremely bright with $L$ from $10^4$ to $10^6 L_\odot$. Nucleosynthesis taking place in the interior of a star has as a consequence that, as evolution proceeds, stellar atmosphere acquires abundance patterns that strongly differ from the chemical composition of the natal cloud in which the star was born. This is referred to as self-pollution: dredge-up episodes bring material from the interior to the surface. Thus the stellar atmosphere becomes enriched or depleted in different chemical elements, e.g., He, C and N, and s-process elements \citep[e.g.][]{1992eatc.conf...92L, 1993ApJ...413..641V, 2005ARA&A..43..435H}. These abundance peculiarities are extreme on the AGB phase. The newly synthesised elements are returned into the ISM, as stars lose mass through winds. Giant stars are thus the primary producers of chemical elements, and much can be learned about stellar evolution from the analysis of abundance ratios as a function of the evolutionary stage of a star. Giant stars serve as a primary gauge of cosmic abundances. They are the key targets in studies of Galactic chemical evolution and stellar archeology. Ongoing observational programmes search for very- and ultra-metal poor (UMP) stars in the Galactic halo and in the bulge \citep{2005ARA&A..43..531B}; many of such objects are old red giants. Some large-scale stellar surveys, such as APOGEE (Apache Point Galactic Experiment from Sloan Digital Sky Survey) focus entirely on RGB stars, for their intrinsic brightness allows to probe very large range of distances in the Galaxy. Individual giant stars can be observed with modern telescopes in the nearby galaxies of the Local Group \citep[e.g.][]{2006AJ....131.2497G, 2010ApJS..191..352K}. Moreover, RGB and AGB stars dominate the integrated light of spatially unresolved systems, such as extra-Galactic globular clusters, dSph, and elliptical galaxies, thus allowing us to determine chemical composition of stellar populations from their composite spectra. The atmospheres of giant stars are not fully in hydrostatic and thermodynamic equilibrium. Theory can explain observations only if very complex physical phenomena are included in the models: outflows, shocks, winds, pulsations, indicating that we are dealing with stars far more complex than the Sun. Violent convective motions penetrate their atmospheres and influence the physical state of matter. The extremely low densities of giant photospheres are to be blamed for the fact that matter and radiation are not in equilibrium. All these phenomena will be briefly reviewed below \citep[for a status update see also the review by][]{2010AN....331..433R}. First, we recap the most recent results from imaging and spectroscopy of cool stars (Sec. \ref{sec:2}), which convincingly show that the stars are more sophisticated compared to their evolutionary predecessors, un-evolved main-sequence stars, and give a glimpse to their complex physics. Sec. \ref{sec:3} describes the methods used for stellar parameter analysis. In Sec. \ref{sec:4}, we review the key results of state-of-the-art modelling of cool giant spectra\footnote{A more detailed discussion of model atmospheres and synthetic spectra for giants are given in the review lecture on 3D NLTE spectroscopy.}.
\label{sec:8} To conclude, modelling spectra of cool giants and supergiants is, doubtlessly, a very challenging problem. Different physical phenomena must be included in the atmosphere and line formation models to allow for an accurate determination of surface parameters: molecular line opacities, deviations from 1D hydrostatic equilibrium and from local thermodynamic equilibrium, chromospheres, winds, and mass loss. We have attained the necessary minimum level of complexity to address the simpler problems, like the determination of metallicity and bolometric flux. However, more detailed investigations, like the analysis of TiO or CO molecular bands and the interpretation of strong lines like H$_\alpha$, Ca H $\&$ K, Mg II, which contain the essential physics, shall await for more complex models including the above-mentioned effects. In the view of the enormous perspectives offered by observing giants and RSGs throughout the Milky Way and in other galaxies, the need for more physically-realistic models is well-justified. \begin{acknowledgement} Figures from the following papers have been reproduced with permission from the authors and from the publishers (c) ESO: Chiavassa et al. A \&A, 535, A22, 2011; Decin et al. A\&A, 548, A113, 2012; Georgy et al. A\&A, 558, A103, 2013; Lancon \& Wood, A\&AS, 146, 217, 2000; Kervella et al., A\&A, 504, 115, 2009. Figures from the following papers have been reproduced by permission of the AAS: Davies et al., ApJ, 767:3, 2013. This work was partly supported by the European Union FP7 programme through ERC grant number 320360. \end{acknowledgement}
14
3
1403.3087
Cool red giants and supergiants are among the most complex and fascinating stars in the Universe. They are bright and large, and thus can be observed to enormous distances allowing us to study the properties of their host galaxies, such as dynamics and chemical abundances. This review lecture addresses various problems related to observations and modelling spectra of red giants and supergiants.
false
[ "chemical abundances", "Universe", "red giants", "Cool red giants", "modelling spectra", "supergiants", "dynamics", "enormous distances", "various problems", "observations", "the Universe", "their host galaxies", "the most complex and fascinating stars", "the properties", "This review lecture", "They", "us" ]
8.798443
10.356273
155
483103
[ "Casagrande, L.", "Silva Aguirre, V.", "Stello, D.", "Huber, D.", "Serenelli, A. M.", "Cassisi, S.", "Dotter, A.", "Milone, A. P.", "Hodgkin, S.", "Marino, A. F.", "Lund, M. N.", "Pietrinferni, A.", "Asplund, M.", "Feltzing, S.", "Flynn, C.", "Grundahl, F.", "Nissen, P. E.", "Schönrich, R.", "Schlesinger, K. J.", "Wang, W." ]
2014ApJ...787..110C
[ "Strömgren Survey for Asteroseismology and Galactic Archaeology: Let the SAGA Begin" ]
102
[ "Research School of Astronomy and Astrophysics, Mount Stromlo Observatory, The Australian National University, Canberra, ACT 2611, Australia; Stromlo Fellow.;", "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, Sydney, NSW 2006, Australia", "NASA Ames Research Center, Moffett Field, CA 94035, USA; SETI Institute, 189 Bernardo Avenue, Mountain View, CA 94043, USA", "Institute of Space Sciences (IEEC-CSIC), Campus UAB, Fac. Ciéncies, Torre C5 parell 2, E-08193 Bellaterra, Spain", "INAF-Osservatorio Astronomico di Collurania, via Maggini, I-64100 Teramo, Italy", "Research School of Astronomy and Astrophysics, Mount Stromlo Observatory, The Australian National University, Canberra, ACT 2611, Australia", "Research School of Astronomy and Astrophysics, Mount Stromlo Observatory, The Australian National University, Canberra, ACT 2611, Australia", "Institute of Astronomy, Madingley Road, Cambridge CB3 0HA, UK", "Research School of Astronomy and Astrophysics, Mount Stromlo Observatory, The Australian National University, Canberra, ACT 2611, Australia", "Stellar Astrophysics Centre, Department of Physics and Astronomy, Aarhus University, Ny Munkegade 120, DK-8000 Aarhus C, Denmark", "INAF-Osservatorio Astronomico di Collurania, via Maggini, I-64100 Teramo, Italy", "Research School of Astronomy and Astrophysics, Mount Stromlo Observatory, The Australian National University, Canberra, ACT 2611, Australia", "Lund Observatory, Department of Astronomy and Theoretical Physics, P.O. Box 43, SE-22100 Lund, Sweden", "Centre for Astrophysics and Supercomputing, Swinburne University of Technology, Hawthorn, VIC 3122, Australia", "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", "Department of Astronomy, The Ohio State University, 140 West 18th Avenue, Columbus, OH 43210-1173, USA; Rudolf-Peierls Centre for Theoretical Physics, University of Oxford, 1 Keble Road, Oxford OX1 3NP, UK; Hubble Fellow.", "Research School of Astronomy and Astrophysics, Mount Stromlo Observatory, The Australian National University, Canberra, ACT 2611, Australia", "National Astronomical Observatories, Chinese Academy of Sciences, Beijing 100012, China" ]
[ "2014A&A...570A..54L", "2014ApJ...790..127B", "2014ApJS..215...19P", "2014MNRAS.444..392C", "2014MNRAS.445.2758R", "2015AJ....149..186P", "2015AJ....150..148H", "2015ASSL..408..169H", "2015ASSP...39...61C", "2015EAS....73..309C", "2015EPJWC.10103001C", "2015IAUS..307..154A", "2015JKAS...48...93G", "2015MNRAS.449.2604D", "2015MNRAS.451.2230M", "2015MNRAS.451.3011C", "2015MNRAS.452.2127S", "2015RAA....15.1125C", "2016A&A...588A..82B", "2016A&A...592A..14R", "2016A&A...594A..39F", "2016A&A...594A.120B", "2016AN....337..805S", "2016AN....337..823S", "2016AN....337..889C", "2016ApJ...822...15S", "2016ApJ...832..133S", "2016ApJS..224....2H", "2016IAUFM..29B.680C", "2016IAUS..319..113C", "2016JKAS...49....9G", "2016MNRAS.455..987C", "2016MNRAS.456.3655M", "2016MNRAS.457.4454T", "2016MNRAS.459..487Y", "2016MNRAS.460.3179W", "2016MNRAS.461.4206V", "2016MNRAS.463.2600L", "2016MmSAI..87..332C", "2016PASP..128l4204L", "2016arXiv160407442H", "2017A&A...597A..30A", "2017A&A...601A..67C", "2017A&A...603A..37G", "2017AJ....153...75K", "2017ApJ...835..172L", "2017ApJ...835..173S", "2017ApJ...837...47V", "2017ApJ...840...77C", "2017ApJ...844..102H", "2017ApJS..229...30M", "2017ApJS..233...23S", "2017EPJWC.16004002C", "2017MNRAS.472..979H", "2017PASP..129k4201K", "2017arXiv170408878G", "2018A&A...609A..58V", "2018A&A...618A.102M", "2018A&A...619A.158V", "2018A&A...620A..76A", "2018A&A...620A.168V", "2018AN....339..134D", "2018ApJ...856..125H", "2018ApJ...859...64H", "2018ApJ...863L..33M", "2018ApJS..237...17S", "2018ApJS..237...21M", "2018IAUS..330..206C", "2018MNRAS.473.2622J", "2018MNRAS.475.5487S", "2018MNRAS.476.5216D", "2018MNRAS.477.5052N", "2018MNRAS.479.4416C", "2019A&A...622A.111S", "2019A&A...623A..60S", "2019A&A...627A.173V", "2019AJ....158..163D", "2019BAAS...51c.488H", "2019MNRAS.482.2770C", "2019MNRAS.485..560C", "2019MNRAS.486.3569H", "2019MNRAS.490.4465R", "2019MNRAS.490.5335S", "2020A&A...635A.164S", "2020ApJ...894...67D", "2021A&A...645A..85M", "2021A&A...656A.155L", "2021AJ....161..170S", "2021ApJ...907...68H", "2021FrASS...7..102J", "2021MNRAS.500...54N", "2021MNRAS.501.5309H", "2021NewA...8401522K", "2021RAA....21..226J", "2022A&A...664A..78F", "2022MNRAS.510..433L", "2022MNRAS.510.1733A", "2023A&A...671A..78M", "2023ApJ...945...20J", "2023MNRAS.523.1089A", "2024ApJS..270...32Q", "2024arXiv240515919L" ]
[ "astronomy" ]
11
[ "Galaxy: stellar content", "stars: abundances", "stars: distances", "stars: fundamental parameters", "stars: oscillations", "surveys", "techniques: photometric", "Astrophysics - Solar and Stellar Astrophysics" ]
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[ "10.1088/0004-637X/787/2/110", "10.48550/arXiv.1403.2754" ]
1403
1403.2754_arXiv.txt
\label{sec:intro} The study of a sizeable number of individual stars in the Milky Way enables us to directly access different phases of its formation and evolution, in a fashion which is still hardly achievable for external galaxies. For obvious reasons, stars in the vicinity of the Sun have been preferred targets for this purpose, starting from the very first \citep[e.g.,][] {gliese57,wallerstein62,vandenBergh62,els62} until the most recent photometric and spectroscopic investigations of the Milky Way, for which {\it Hipparcos} astrometric distances are also available out to $\simeq 100$~pc \citep[e.g.,][]{nordstrom04,reddy06,fb08,c11,ad13,b13}. Properties of stars in the solar neighborhood, in particular ages and metallicities, are still the main constraint for Galactic chemo(dynamical) models \citep[e.g.,][]{cescutti07,ralph09a,kn11,m12,wz13} and provide important clues to understanding the main processes at play in galaxy formation and evolution \citep[e.g.,][]{pilk12,bird13}. This sort of study is now fostered by a number of spectroscopic and photometric surveys, targeting different and fainter components of the Milky Way (e.g.,~``RAVE'' \citealt{steinmetz06}; ``SDSS-SEGUE'' \citealt{ya09}; ``SDSS-APOGEE'' \citealt{ahn13}; ``VVV'' \citealt{minniti10}; ``Gaia-ESO'' \citealt{gge}; ``GALAH'' \citealt{ken}), although astrometric distances for the targets in these surveys will not be available until the {\it Gaia} spacecraft releases its data \citep{lp96:gaia,p01:gaia}. A common feature of all past and current stellar surveys is that while it is relatively straightforward to derive some sort of information on the chemical composition of the targets observed (and in many cases even detailed abundances), that is not the case when it comes to masses, radii, distances and, in particular, ages. Even when accurate astrometric distances are available to allow comparison of stars with isochrones, the derived ages are still highly uncertain, and statistical techniques are required to avoid biases \citep[e.g.,][]{pont04,Jorgensen05}. Furthermore, isochrone dating is meaningful only for stars in the turnoff and subgiant phase, where stars of different ages are clearly separated on the HR diagram. This is in contrast, for example, to stars on the red giant branch, where isochrones with vastly different ages can fit equally well observational constraints such as effective temperatures, metallicities and surface gravities within their errors \citep[e.g.,][for a review]{soderblom}. Thus, alternative ways to precisely determine masses and radii of stars are the only way forward. By measuring oscillation frequencies in stars, asteroseismology allows us to determine fundamental physical quantities, in particular masses and radii, which otherwise would be inaccessible in single field stars, and which can be used to obtain information on stellar distances and ages \citep[e.g.,][for a review]{cm13}. Individual frequencies can yield an accuracy of just a few percent on those parameters but their exploitation is more demanding, both observationally and theoretically \citep{vsa13}. Global asteroseismic observables (see Section \ref{sec:astero}) on the contrary are easier to detect and analyze, yet able to provide the aforementioned parameters for a large number of stars with an accuracy that is still much better than achievable by isochrone fitting in the traditional sense. Thanks to space-borne missions such as {\it CoRoT} \citep{corot06} and {\it Kepler} \citep{gilli10} average oscillation frequencies are now robustly detected in more than 500 main-sequence and subgiant stars, and over $13,000$ giants \citep[e.g.,][]{derid09,verner11,stello13}. Asteroseismology is thus emerging as a new tool for studying stellar populations, and initial investigations in this direction have already been done \citep[][]{chap11,miglio13}. However, until now asteroseismic studies of stellar populations had only coarse information on ``classical '' stellar parameters such as $\teff$ and $\feh$. Coupling classical parameters with seismic information not only improves the seismic masses and ages obtained for stars \citep[][]{lm09,chap13}, but also allows us to address important questions in stellar and Galactic evolution. To fully harvest the potential that asteroseismology brings to these studies, classical stellar parameters are vital, yet unavailable for a large sample of stars with detected oscillations. The main purpose of the Kepler Input Catalog \citep[KIC,][]{kic} was to separate dwarfs from giants, and it is therefore inadequate for high precision stellar and Galactic studies due to significant biases in its stellar parameters \citep[e.g.,][]{mz11,pi12,thy12}. APOGEE will eventually provide parameters for thousands of {\it Kepler} giants, and similarly the Gaia-ESO Survey and GALAH for the {\it CoRoT} fields. The advantages of the spectroscopic surveys are obvious, yet photometry still offers a powerful and complementary alternative. Here, we present the first results from the ongoing Str\"omgren survey for Asteroseismology and Galactic Archaeology (SAGA) in the {\it Kepler} field. The goals of our photometric survey are manifold. First, the Str\"omgren system was envisaged with the ultimate purpose of studying Galactic stellar populations, and designed to provide reliable stellar parameters, in particular metallicities \citep[][]{stromgren63,stro87}. Thus, even compared to multi-fiber spectroscopy, wide-field Str\"omgren imaging is very efficient. It also has the advantage that no pre-selection is made on targets: all stars that fall in a given brightness regime across the instrument field of view will be observed. This greatly simplifies the selection function, which in our case is essentially dictated by the {\it Kepler} satellite \citep[e.g.,][]{bata,fk13}. Further, relatively faint magnitudes can be probed at high precision even on a small-class telescope, making it possible to readily derive metallicities over a large magnitude range. The {\it Kepler} seismic sample presented in this work can thus be used as reference e.g.~for assessing the accuracy at which stellar parameters can be obtained for stars having only photometric measurements, such as the planet host stars we observe in SAGA. At the same time, the sensitivity of Str\"omgren colors to metallicity, coupled with seismic ages, help immensely to calibrate other metallicity and age dating techniques, thus creating powerful new tools to study more remote Galactic populations than previously possible. Bearing in mind the expected precision at which {\it Gaia} will deliver astrophysical parameters for stars of different brightness \citep{liu12}, the above science case highlight the importance of having an all-sky Str\"omgren survey going fainter than any of the current large spectroscopic surveys \citep{wang}. The SAGA survey also represents a natural extension of the Geneva Copenhagen Survey \citep[GCS,][]{nordstrom04}, the only all-sky Str\"omgren survey currently available, and a benchmark for Galactic studies. Similar to our latest revision of the GCS \citep{c11}, we combine Str\"omgren metallicities with broad-band photometry to obtain effective temperatures ($\teff$) and metallicities ($\feh$) for all our targets via the Infrared Flux Method (IRFM). This facilitates the task of placing SAGA and GCS on the same scale. However, there are marked differences between the two surveys: the GCS is an all-sky, shallow survey limited to main-sequence and subgiant stars closer than $\simeq 100$~pc ($40$~pc volume limited). The {\it Kepler} targets observed by SAGA are primarily giants located between $\simeq1$ and $\simeq6$~kpc in a specific region of the Galactic disk (Figures \ref{f:fov} and \ref{f:trigo}). The use of giants as probes of Galactic Archaeology is possible since it is relatively straightforward to derive ages for these stars once classical stellar parameters are coupled with asteroseismology. This was not the case for the GCS, where isochrone fitting was used, and thus limited to main-sequence and subgiant stars with known astrometric distances. On the other hand, stars in the GCS have kinematic information, which is not available for the SAGA targets. The different distance ranges sampled by the GCS and SAGA makes them complementary: the stellar properties measured within the solar neighborhood in the former survey can be dynamically stretched across several kpc using kinematics. In contrast, the larger distance range sampled by the giants in SAGA provides {\it in situ} measurements of various stellar properties over $\simeq5$~kpc. In this paper we present the observing strategy and data reduction of the Str\"omgren survey. After cross-matching this dataset with stars in the {\it Kepler} field with seismic information, we derive classical and seismic stellar parameters for almost 1,000 targets. Our approach is characterized by the powerful combination of classical and seismic parameters, to our knowledge the first of this kind, and a careful treatment of the errors, allowing us to derive self-consistent effective temperatures, distances, masses and radii with a typical precision of a few percent, in addition to metallicities as we discuss in more detail later. With these parameters in hand, stellar ages are a straightforward by-product, although more caution must be used: we leave their estimation to a subsequent publication.
In this paper we have presented Str\"omgren photometry of a stripe in the {\it Kepler} field. While the ultimate purpose is to use these observations to provide homogeneous and reliable stellar parameters for both candidate planet host and seismic stars, the geometry chosen for our observations already enables Galactic studies using stars with detected oscillations. The strength of our approach is in the coupling of classical and asteroseismic parameters: effective temperatures (from the InfraRed Flux Method), metallicities (from Str\"omgren indices), masses, and radii (from seismology) are derived iteratively and self-consistently, thus giving access to other quantities such as reddening and distances, and providing a complete picture of the seismic population in our observed Galactic fields. The open cluster NGC\,6819 is located at the base of the observed stripe and offers an important benchmark to verify the soundness of our results. In particular, it allows us to anchor the metallicity scale of giants on the same zero-point used for investigating properties in the solar neighborhood via dwarfs and subgiants from the Geneva-Copenhagen Survey \citep{c11}. This is crucial if for example we wish to use nearby stars and/or giants (up to $\simeq 6$~kpc across the Galactic disk in this work) for the purpose of understanding the formation and evolution of stellar populations in the Milky Way. The sample presented here covers latitudes between $8$ and $20$ degrees above the Galactic plane. The stellar parameters derived here are already well suited for a number of investigations; in particular the seismic classification allows us to distinguish between stars ascending the red giants branch burning hydrogen in a shell, and those which have also ignited helium burning in their core. Concerning the latter, the precision achieved allows us to discern between primary and secondary clump, hence whether the core ignites degenerately or not. With this information, and the parameters so far derived, we are thus in the position of using these stars to compute reliable stellar ages and investigate, for example the age-metallicity relation and the age/metallicity gradients across this part of the Galactic disk. Further, calibrating photometric metallicities and age dating techniques on the present sample, a deep all-sky Str\"omgren survey promises a leading role for Galactic studies.
14
3
1403.2754
Asteroseismology has the capability of precisely determining stellar properties that would otherwise be inaccessible, such as radii, masses, and thus ages of stars. When coupling this information with classical determinations of stellar parameters, such as metallicities, effective temperatures, and angular diameters, powerful new diagnostics for Galactic studies can be obtained. The ongoing Strömgren survey for Asteroseismology and Galactic Archaeology has the goal of transforming the Kepler field into a new benchmark for Galactic studies, similar to the solar neighborhood. Here we present the first results from a stripe centered at a Galactic longitude of 74° and covering latitude from about 8° to 20°, which includes almost 1000 K giants with seismic information and the benchmark open cluster NGC 6819. We describe the coupling of classical and seismic parameters, the accuracy as well as the caveats of the derived effective temperatures, metallicities, distances, surface gravities, masses, and radii. Confidence in the achieved precision is corroborated by the detection of the first and secondary clumps in a population of field stars with a ratio of 2 to 1 and by the negligible scatter in the seismic distances among NGC 6819 member stars. An assessment of the reliability of stellar parameters in the Kepler Input Catalog is also performed, and the impact of our results for population studies in the Milky Way is discussed, along with the importance of an all-sky Strömgren survey. <P />Based on observations made with the Isaac Newton Telescope operated on the island of La Palma by the Isaac Newton Group in the Spanish Observatorio del Roque de los Muchachos of the Instituto de Astrofísica de Canarias.
false
[ "field stars", "stars", "seismic information", "Galactic studies", "NGC", "NGC 6819 member stars", "Galactic Archaeology", "distances", "Galactic", "population studies", "effective temperatures", "powerful new diagnostics", "masses", "stellar parameters", "radii", "surface gravities", "the negligible scatter", "the seismic distances", "classical and seismic parameters", "stellar properties" ]
7.587003
11.628599
81
693828
[ "Garriga, Jaume", "Urakawa, Yuko" ]
2014JHEP...06..086G
[ "Holographic inflation and the conservation of ζ" ]
23
[ "Departament de Física Fonamental i Institut de Ciències del Cosmos, Universitat de Barcelona", "Departament de Física Fonamental i Institut de Ciències del Cosmos, Universitat de Barcelona; Department of Physics and Astrophysics, Nagoya University" ]
[ "2014JHEP...06..052K", "2014arXiv1405.1762L", "2015JCAP...01..028G", "2015JCAP...01..033M", "2015PhRvD..91f4018W", "2016JCAP...02..036G", "2016JCAP...06..020T", "2016JCAP...10..030G", "2016JHEP...12..028I", "2017JCAP...03..060B", "2017PhRvD..95l3505A", "2017PhRvD..96f5011I", "2017PhRvL.118d1301A", "2018JCAP...01..024C", "2018JCAP...07..021B", "2019JCAP...09..044C", "2020JHEP...12..026N", "2020PhRvD.101b1901N", "2021JHEP...04..166I", "2022arXiv220406342V", "2023Univ....9...59S", "2023arXiv230814690C", "2024JHEP...05..056C" ]
[ "astronomy", "physics" ]
0
[ "Gauge-gravity correspondence", "Renormalization Group", "High Energy Physics - Theory", "Astrophysics - Cosmology and Nongalactic Astrophysics", "General Relativity and Quantum Cosmology" ]
[ "1969PhRv..179.1499W", "1998AdTMP...2..231M", "1998AdTMP...2..253W", "1998PhLB..428..105G", "2000JHEP...08..003D", "2000PhRvD..62d3527W", "2001CMaPh.217..595D", "2001JHEP...10..034S", "2001JHEP...11..049S", "2002NuPhB.631..159B", "2002PhRvD..65j4039B", "2003JHEP...05..013M", "2003JHEP...07..051L", "2003PhRvD..67l3504W", "2004CQGra..21L..65M", "2004JHEP...01..070V", "2004JHEP...07..062L", "2005JCAP...05..004L", "2005PhRvD..72j3501L", "2006JCAP...06..001S", "2009JCAP...01..021G", "2009JCAP...11..020G", "2010JPhCS.222a2007M", "2010PhRvD..81b1301M", "2011CQGra..28g2001N", "2011JCAP...05..013M", "2011JCAP...06..030M", "2011JCAP...06..032V", "2011JCAP...10..021G", "2011JHEP...10..038K", "2011PThPh.125.1067U", "2011arXiv1104.2621H", "2011arXiv1108.5735A", "2012PhRvD..86h5025S", "2012mgm..conf.2315M", "2013CQGra..30w3001T", "2013JCAP...03..005S", "2013JCAP...07..033G", "2013JCAP...11..011K", "2013JHEP...02..151A", "2013JHEP...04..047B", "2013JHEP...07..015M", "2013JHEP...09..148S", "2013JHEP...10..071M", "2013PTEP.2013f3E02T", "2013arXiv1306.3999B", "2013arXiv1311.0755B", "2014A&A...571A..22P", "2014ForPh..62..389K", "2014JHEP...02..124P", "2014JHEP...07..011G", "2014JHEP...08..081M", "2014PhRvL.112x1101B", "2014arXiv1405.1762L" ]
[ "10.1007/JHEP06(2014)086", "10.48550/arXiv.1403.5497" ]
1403
1403.5497_arXiv.txt
\label{Sec:Introduction} Inflation has become a standard paradigm for describing the origin of cosmological perturbations~\cite{Ade:2013rta, Ade:2014xna}. In fact, current observational data is in good agreement with single field models, with just one inflaton field~\cite{Ade:2013rta}. On the other hand, it has been suggested that inflation may be described holographically by means of a dual field theory at the future boundary. According to the gauge/gravity correspondence, the strongly (weakly) coupled phase of bulk gravity corresponds to the weakly (strongly) coupled phase of the dual boundary theory. Because of that, holography may open up new insights on the study of the very early universe, near the Planck scale, where non-perturbative gravitational effects may play a role. The gauge/gravity duality was initially advocated for asymptotically anti-de Sitter (AdS) space times~\cite{Maldacena1997, GKP, Witten1998}. Such duality cannot be immediately applied to inflationary cosmology, where the spacetime is similar to de Sitter (dS) rather than AdS. Nonetheless, by analogy, there have been several suggestions that a $(d+1)$-dimensional inflationary evolution may be dual to a quantum field theory (QFT) on a $d$-dimensional space with Euclidean signature. Following the work by Strominger~\cite{Strominger, Strominger2} and Witten~\cite{Witten}, this possibility has been further investigated in the context of dS~\cite{Bousso:2001mw, Harlow:2011ke, Anninos:2011ui} and quasi-dS spacetimes~\cite{Strominger2, Maldacena02}. In Refs.~\cite{Maldacena02, Seery:2006tq, vdS, LM03, LM04, Shiu, Mata:2012bx, Ghosh:2014kba, JYsingle, Larsen:2014wpa}, the duality was discussed by including cosmological perturbations. (See also Refs.~\cite{Banks:2013qra, Banks:2013qpa} and Ref.~\cite{Pimentel:2013gza}.) The holographic description of inflation has also been studied by using the so-called domain wall/cosmology correspondence, where cosmological solutions are constructed by analytically continuing from domain wall solutions~ \cite{MS_HC09, MS_HC10, MS_HCob10, MS_NG, MS_NGGW, BMS} (see also Ref.~\cite{Kiritsis:2013gia}). The implementation of the duality requires a concrete dictionary, relating cosmological observables in the bulk with field theory observables at the boundary. However, this relation hasn't been clearly understood, except perhaps in certain limits, such as the vicinity of a dS fixed point. In particular, it is not clear which cosmological variable corresponds to the renormalization scale $\mu$. In Refs.~\cite{BMS, AJ08, AJ09, Alex11} , it was argued that for the case of dS spacetime, $\mu$ should be proportional to the scale factor, $\mu \propto a$, but the relations suggested in these references differ from each other when the solutions deviate from dS spacetime. One may expect that in quasi-dS spacetimes the cosmological evolution in the bulk will be still described by the renormalization group (RG) flow in the boundary. The purpose of this paper is to examine this naive expectation, by computing the evolution of the primordial curvature perturbation $\zeta$. This plays an important role because in standard cosmological perturbation theory (CPT) $\zeta$ is generically conserved for adiabatic perturbations at large scales (this will be reviewed more precisely in Section 5). If the time evolution in the bulk consistently corresponds to the RG flow in the dual boundary QFT, the correlators of $\zeta$ predicted in the boundary QFT should be independent of $\mu$ in the large scale limit. In this paper, we examine whether the RG flow in the boundary QFT predicts that the correlators of $\zeta$ become $\mu$ independent or not. The outline of this paper is as follows. In Sec.~\ref{Sec:Preliminaries}, after we describe our setup, following Ref.~\cite{JYsingle}, we provide a way to calculate the correlators of $\zeta$ by using the wave function of the bulk spacetime. To consider the boundary QFT which is dual to the single field inflation, we introduce a single deformation term to the boundary action, which lets the QFT deviate from the conformal field theory (CFT). In Sec.~\ref{Sec:dCFT}, we discuss the solution of the RG equation in the boundary theory. In Sec.~\ref{Sec:vertexfn}, using the gauge transformation, we derive the relation between the correlators of $\zeta$ and the correlators of the boundary operator. Then, in Sec.~\ref{Sec:conservation}, computing the correlators of $\zeta$, we investigate their $\mu$ dependence at large scales.
In this paper, we examined whether the curvature perturbation $\zeta$ is conserved under the change of $\mu$ in a local boundary theory with the action (\ref{Exp:Sren}). This corresponds to a generic CFT with a single deformation operator. In order to relate the boundary correlators to the correlators of the gauge invariant curvature perturbation $\zeta$, we have assumed that $\zeta$ is locally related to the scalar field perturbation in the flat gauge, $\delta\phi_f$ (or equivalently the perturbation of the coupling $\delta g_f$) as in Eq.~(\ref{ASM:local}). This assumption certainly holds in standard cosmological perturbation theory at large scales. Also, we have assumed that the boundary operator ${\cal O}$ is renormalized multiplicatively as in Eq.~(\ref{ASM:ren}). Solving the RG flow, we found that the power spectrum of $\zeta$ is conserved if we identify the renormalization scale $\mu$ with the scale factor $a$ in the bulk, as in Eq.~(\ref{Rel:mua}). But then, it follows that the bi-spectrum of $\zeta$ cannot be conserved along the entire RG flow. The only exception is the particular case where the scaling dimension is constant, as given in Eq.~(\ref{tempbeta}). This special case leads to an exact power law spectrum of the form (\ref{Eq:Pkconserved}). In order to have conservation of $\zeta$ along a generic RG flow, we need to abandon at least one of the following three assumptions: the local relation Eq.~(\ref{ASM:local}), the multiplicative renormalization Eq.~(\ref{ASM:ren}) or the assumption of a local boundary theory with the single deformation operator. The local relation Eq.~(\ref{ASM:local}) follows from Eqs.~(\ref{Rel:gphi}) and (\ref{Rel:zetadphi}). Instead of imposing Eq.~(\ref{Rel:gphi}), one may need to seek for a more non-trivial relation between $g$ and $\phi$ (this issue has been discussed in AdS/CFT. See, for instance, Ref.~\cite{Megias:2014iwa}). A simple generalization to the local function $g(\mu,\, \bm{x})=g[\phi(t(\mu),\, \bm{x})]$ is just a change of variable describing the field, and will not help to preserve the conservation. The second and third assumptions are concerned with renormalization. In this paper, we assumed that the boundary QFT can be renormalized by introducing the counterterm and the wave function renormalization as in Eqs.~(\ref{Exp:Sren}) and (\ref{ASM:ren}), respectively. Nevertheless, in general, one may need to introduce more than one deformation operator to perform the renormalization. In addition, the QFT may become non-local after the renormalization. These cases were not addressed in the present paper. When the boundary theory contains more than one deformation operator, the corresponding cosmological evolution will be governed by several scalar fields. Notice, however, that in this case the standard CPT does not predict the conservation of $\zeta$ any longer. The possibility of generalizing the duality to the case of a non-local boundary theory is at present not very well understood. In particular, it is not clear whether the locality (non-locality) in one side of the duality implies the locality (non-locality) on the other side. A relevant discussion can be found in Ref.~\cite{Sundrum:2011ic}, but to our knowledge this issue has not been fully resolved, at least in the case when the deviation from dS spacetime becomes important (even in the large $N$ limit). If a non-local boundary theory can be dual to a local bulk theory with a single field, the conservation of $\zeta$ should be predicted also from the boundary computation. We leave this issue for future research. In this paper, we discussed the asymptotically dS spacetimes and the dual boundary QFT. Since the asymptotically dS spacetimes which we have considered can be transformed into the asymptotically AdS spacetime by analytic continuation, the analog of $\zeta$ will be conserved along the holographic direction also in the case of asymptotically AdS spacetimes. Possible implications of our results in this broader context are currently under investigation.
14
3
1403.5497
In a holographic description of inflation, cosmological time evolution in the bulk is expected to correspond to the renomalization group (RG) flow in a dual boundary theory. Here, we analyze this expectation by computing the correlation functions of the curvature perturbation ζ holographically. For this purpose, we use a deformed conformal field theory at the boundary, with a single deformation operator. In standard single field models of inflation, ζ is known to be conserved at large scales under very general conditions. However, we find that this is not generically the case in the dual description. The requirement that higher correlators of ζ should be conserved severely restricts the possibilities for the RG flow. With such restriction, the power spectrum P <SUB> ζ </SUB> must follow an exact power law, at least within the regime of validity of conformal perturbation theory.
false
[ "conformal perturbation theory", "very general conditions", "standard single field models", "a dual boundary theory", "ζ", "a single deformation operator", "cosmological time evolution", "a deformed conformal field theory", "large scales", "inflation", "the curvature perturbation ζ", "the RG flow", "validity", "an exact power law", "the dual description", "higher correlators", "(RG) flow", "the renomalization group", "such restriction", "the boundary" ]
11.192509
-0.476608
89
746455
[ "Chamandy, Luke", "Shukurov, Anvar", "Subramanian, Kandaswamy", "Stoker, Katherine" ]
2014MNRAS.443.1867C
[ "Non-linear galactic dynamos: a toolbox" ]
40
[ "Inter-University Centre for Astronomy and Astrophysics, Post Bag 4, Ganeshkhind, Pune 411007, India", "Inter-University Centre for Astronomy and Astrophysics, Post Bag 4, Ganeshkhind, Pune 411007, India; School of Mathematics and Statistics, Newcastle University, Newcastle upon Tyne NE1 7RU, UK", "Inter-University Centre for Astronomy and Astrophysics, Post Bag 4, Ganeshkhind, Pune 411007, India", "School of Mathematics and Statistics, Newcastle University, Newcastle upon Tyne NE1 7RU, UK" ]
[ "2013pss5.book..641B", "2014ApJS..215....1V", "2015A&A...578A..94M", "2015AN....336..991B", "2015ApJ...799...35V", "2015ApJ...808...28C", "2015MNRAS.446L...6C", "2015MNRAS.450.3472R", "2016A&A...588A..18M", "2016ASSL..435..585D", "2016ApJ...817...12P", "2016ApJ...818...28S", "2016ApJ...818L..10T", "2016ApJ...833...43C", "2016JApA...37...42R", "2016JFM...798..696S", "2016MNRAS.462.4402C", "2017ApJ...841...67F", "2017MNRAS.468.3657C", "2017MNRAS.469.1466Z", "2017MNRAS.469.4806H", "2018JCAP...08..049B", "2018JPlPh..84c7302Z", "2018MNRAS.473.4544S", "2018MNRAS.481.1300C", "2018arXiv181007246C", "2019A&A...623A.113S", "2019Galax...7...91B", "2019Galax...8....4B", "2019MNRAS.483.2424R", "2019MNRAS.487.1498W", "2020MNRAS.499.2076S", "2020MNRAS.499.3673W", "2021RAA....21...56M", "2021arXiv211201193C", "2024ApJ...961....7G", "2024ApJ...966...78C", "2024MNRAS.530..382R", "2024MNRAS.tmp.1483J", "2024arXiv240104015K" ]
[ "astronomy" ]
7
[ "dynamo", "magnetic fields", "MHD", "galaxies: magnetic fields", "galaxies: spiral", "galaxies: structure", "Astrophysics - Astrophysics of Galaxies" ]
[ "1976JFM....77..321P", "1979Ap&SS..66..369R", "1992ApJ...393..165V", "1993MNRAS.265..649S", "1994PhRvL..72.1651G", "1995A&A...297..159K", "1995MNRAS.275..191M", "1996ARA&A..34..155B", "1998MNRAS.297..860M", "1998MNRAS.299L..21S", "1999A&A...343..120M", "1999A&A...350..423R", "2000A&A...358.1142M", "2000A&A...361L...5K", "2000ApJ...534..984B", "2000MNRAS.318..925F", "2001A&A...370..635B", "2001ApJ...550..752V", "2001GApFD..94..135P", "2002A&A...387..453K", "2003GApFD..97..249R", "2003and..book..269B", "2005LNP...664..113S", "2005PhR...417....1B", "2005mpge.conf..117E", "2006A&A...448L..33S", "2006ApJ...648L..71S", "2007A&A...470..539B", "2007MNRAS.377..874S", "2008A&A...490.1005T", "2008ApJ...677L..17C", "2008RPPh...71d6901K", "2009MNRAS.398.1414B", "2010AN....331..130M", "2010ASPC..438..197F", "2012ApJ...754L..35H", "2013A&A...556A.147M", "2013ApJ...772L..28B", "2013MNRAS.428.3569C", "2013MNRAS.429..967G", "2013MNRAS.433.3274C", "2014ApJ...780..144V", "2014GApFD.108..568J", "2014MNRAS.437..562C", "2014PhRvL.112l5003E", "2015SSRv..188...59B" ]
[ "10.1093/mnras/stu1274", "10.48550/arXiv.1403.2562" ]
1403
1403.2562_arXiv.txt
The mean-field dynamo theory provides an appealing explanation of the presence and structure of large-scale magnetic fields in disc galaxies \citep{Ruzmaikin+88,Beck+96,Brandenburg+Subramanian05a,Shukurov05,Kulsrud+Zweibel08}. The dynamo time scale is shorter than the galactic lifetime, and the energy densities of the large-scale galactic magnetic fields and interstellar turbulence are observed to be of the same order of magnitude. It is thus plausible that the galactic large-scale dynamos are normally in a non-linear, statistically steady state. Recent progress in dynamo theory has lead to physically motivated non-linear models where the steady state is achieved through the magnetic helicity balance in the dynamo system \citep[reviewed by][]{Brandenburg+Subramanian05a,Blackman14}. To avoid a catastrophic suppression of the mean induction effects of turbulence, the magnetic helicity of random magnetic fields should be removed from the system. In galaxies, this can be achieved through the advection of magnetic fields from the disc to the halo by the galactic fountain or wind \citep{Shukurov+06,Sur+07}, diffusive flux \citep{Kleeorin+00,Kleeorin+02} and helicity flux relying on the anisotropy of the interstellar turbulence \citep{Vishniac+Cho01,Vishniac+Shapovalov14} \citep[see ][for an application to galaxies]{Sur+07}. The first of these mechanisms is the simplest in physical and mathematical terms, and involves galactic parameters that are reasonably well constrained observationally. Most of the earlier analytical and numerical results in the non-linear mean-field disc dynamo theory rely on a much simpler form of non-linearity in the dynamo equations, the so-called algebraic $\alpha$-quenching that is based on a simple, explicit form of the dependence of the dynamo parameters, usually the $\alpha$-coefficient, on the magnetic field. In a thin layer, such as a galactic or accretion disc, this allows one to obtain a wide range of analytical and straightforward numerical solutions using simple and yet accurate approximations \citep[e.g.,][]{Shukurov04}. One of the advantages of the resulting theory of galactic magnetic fields is that all its essential parameters can be expressed in terms of observable quantities (the angular velocity of rotation, thickness of the gas layer, turbulent velocity, etc.). As a result, theory of galactic magnetic fields has been better constrained and verified by direct comparison with observations than, arguably, any other astrophysical dynamo theory. Such comparisons require relatively simple, preferably analytical, approximations to the solutions of the dynamo equations. In this paper we consider numerical solutions of thin-disc dynamo equations with a dynamic non-linearity involving magnetic helicity balance and compare them with a wide range of simpler solutions to develop a set of accessible tools to facilitate applications of the theory. The paper is organized as follows. In Sections~\ref{sec:dynamo}-\ref{sec:equations}, we present the theoretical background and a review of each of the approximations discussed. This is followed by a detailed comparison of the solutions resulting from various physical and mathematical approximations in Section \ref{sec:results}. In particular, in Section~\ref{sec:alg} we provide an in-depth comparison of the dynamical and algebraic non-linearities. Our overall conclusion is that the earlier, simple models, when applied judiciously, reproduce comfortably well solutions with the dynamical non-linearity. Section~\ref{sec:examples} provides examples of applications of the toolbox, namely the magnetic pitch angle problem and the spiral arm-interarm contrasts in magnetic field. We present a summary and general conclusions in Section \ref{sec:conc}. The details of the asymptototic solutions studied, namely the perturbation and no-$z$ solutions, are given in Appendices~\ref{sec:pert_app} and \ref{sec:noz_app}, respectively. Throughout this paper, we use cylindrical polar coordinates $(r,\phi,z)$ with the origin at the disc centre and the $z$-axis aligned with the angular velocity of rotation $\bm\Omega$.
\label{sec:conc}\label{sec:summary} We have discussed various simple approximate approaches to estimate the strength of the large-scale galactic magnetic fields, their pitch angle and dynamo thresholds, and compared them with numerical solutions. In particular, we compared the non-linear states established due to magnetic helicity conservation with those obtained with a much simpler, and easier to analyze, heuristic algebraic form of the dynamo non-linearity. These approaches complement one another. For example the perturbation solution of Section~\ref{sec:pert} provides a reasonably accurate form of the distribution of magnetic field across the disc, whereas the no-$z$ approximation of Section~\ref{sec:noz} gives useful results for variables averaged across the disc. Remarkably, and reassuringly, where they overlap, all of these methods result in similar solutions. Most importantly, results obtained with the dynamical non-linearity that involves advective and diffusive fluxes of magnetic helicity are very much consistent with those from the algebraic $\alpha$-quenching. We suggest how the latter can be modified to achieve quite a detailed agreement. Magnetic lines produced by the mean-field dynamo are believed to be trailing with respect to the galactic rotation because the galactic angular velocity decreases with galactocentric radius. We have found, however, that steady-state magnetic fields obtained for the dynamical non-linearity are trailing near the galactic midplane but leading closer to the disc surface (where $\mean{B}_r$ changes sign) if an outflow is present. This effect is more pronounced when the galactic outflow is stronger or the dynamo number is higher as compared with its critical value. This feature is new and unexpected, as it is not reproduced in models with algebraic quenching. This makes it reasonable to expect that leading magnetic spirals may be observable in the disc-halo interface regions of spiral galaxies (or even higher in the halo). To what extent this feature persists if the boundary conditions are varied or if the galactic halo is included is a question that merits future investigation. It is also useful for applications that marginal kinematic solutions of the dynamo equations in a thin disc (i.e., those that neither grow nor decay) reproduce with high accuracy non-linear steady-state solutions. The simple analytical perturbation solutions of kinematic dynamo equations, here generalized to include magnetic field advection in a galactic outflow, are particularly useful in this respect. It has been shown here and by \citet{Ji+13} that they remain accurate beyond their formal range of applicability and can be used for the range of dynamo numbers $-50\la D\la 0$ typical of galactic discs. Here we have also shown that these solutions can be used as a good approximation to the non-linear states. We have also refined the no-$z$ approximation to allow for vertical advection of the mean magnetic field, as well as advective and diffusive helicity fluxes. We note that advection affects dynamo action through three channels: by reducing the critical dynamo number, by helping the turbulent diffusion to remove flux from the dynamo active region, and by the removal of small-scale magnetic helicity. The heuristic diffusive flux of magnetic helicity has previously been observed in numerical simulations. The models investigated here are somewhat simplified compared to real galaxies. It is worth extending the models to include spatial variation of $\etat$, additional terms in the mean electromotive force, and other contributions to the magnetic helicity flux. The possible importance of $\etat$-quenching, in addition to $\alpha$-quenching \citep{Gressel+13a}, also deserves exploration. More refined modelling will enable better comparison with real galaxies. In summary, much of the earlier work on galactic dynamos modeled the saturation of the dynamo using algebraic quenching of the $\alpha$ effect. We show here that this algebraic quenching non-linearity (which predates dynamical quenching theory but is still widely used in the dynamo literature) is a good approximation to dynamical quenching for the galactic mean-field dynamo. We also extend the standard algebraic quenching formula to make it more accurate. In addition, we suggest three simple tools, namely marginal kinematic solutions, critical asymptotic solutions from perturbation theory, and no-$z$ solutions, and show that all agree remarkably well with the numerical solutions of the non-linear dynamo. Particularly useful are the analytical expressions \eqref{Brpert} and \eqref{Bppert} for the vertical profiles of $\mbr$ and $\mbp$, as well as equations \eqref{Bsat} and \eqref{psat} for the saturation field strength $B$ and pitch angle $p$, which, when used along with the analytical expression \eqref{gammapert} for the kinematic growth rate $\gamma$, comprise a surprisingly efficient guide to the parameter space of galactic dynamos.
14
3
1403.2562
We compare various models and approximations for non-linear mean-field dynamos in disc galaxies to assess their applicability and accuracy, and thus to suggest a set of simple solutions suitable to model the large-scale galactic magnetic fields in various contexts. The dynamo saturation mechanisms considered are the magnetic helicity balance involving helicity fluxes (the dynamical α-quenching) and an algebraic α-quenching. The non-linear solutions are then compared with the marginal kinematic and asymptotic solutions. We also discuss the accuracy of the no-z approximation. Although these tools are very different in the degree of approximation and hence complexity, they all lead to remarkably similar solutions for the mean magnetic field. In particular, we show that the algebraic α-quenching non-linearity can be obtained from a more physical dynamical α-quenching model in the limit of nearly azimuthal magnetic field. This suggests, for instance, that earlier results on galactic disc dynamos based on the simple algebraic non-linearity are likely to be reliable, and that estimates based on simple, even linear models are often a good starting point. We suggest improved no-z and algebraic α-quenching models, and also incorporate galactic outflows into a simple analytical dynamo model to show that the outflow can produce leading magnetic spirals near the disc surface. The simple dynamo models developed are applied to estimate the magnetic pitch angle and the arm-interarm contrast in the saturated magnetic field strength for realistic parameter values.
false
[ "simple solutions", "leading magnetic spirals", "realistic parameter values", "various models", "disc galaxies", "the saturated magnetic field strength", "non-linear mean-field dynamos", "the mean magnetic field", "galactic outflows", "nearly azimuthal magnetic field", "the algebraic α-quenching non", "The non-linear solutions", "various contexts", "helicity fluxes", "a more physical dynamical α-quenching model", "a simple analytical dynamo model", "the large-scale galactic magnetic fields", "the magnetic helicity balance", "The simple dynamo models", "the magnetic pitch angle" ]
12.784369
10.033153
-1
745439
[ "Danehkar, A.", "Todt, H.", "Ercolano, B.", "Kniazev, A. Y." ]
2014MNRAS.439.3605D
[ "Observations and three-dimensional photoionization modelling of the Wolf-Rayet planetary nebula Abell 48" ]
24
[ "Department of Physics and Astronomy, Macquarie University, Sydney, NSW 2109, Australia", "Institut für Physik und Astronomie, Universität Potsdam, Karl-Liebknecht-Str. 24/25, D-14476 Potsdam, Germany", "Universitäts-Sternwarte München, Ludwig-Maxmilians Universität München, Scheinerstr. 1, D-81679 München, Germany; Exzellenzcluster Universe, Technische Universität München, Boltzmannstr. 2, D-85748 Garching, Germany", "South African Astronomical Observatory, PO Box 9, 7935 Observatory, Cape Town, South Africa; Southern African Large Telescope Foundation, PO Box 9, 7935 Observatory, Cape Town, South Africa; Sternberg Astronomical Institute, Lomonosov Moscow State University, 119992 Moscow, Russia" ]
[ "2012msso.prop..158D", "2013aao..prop...19P", "2014MNRAS.441.3065S", "2014PhDT........76D", "2015ApJ...815...35D", "2015MNRAS.449L..56D", "2015PKAS...30..159D", "2015PKAS...30..163D", "2015RMxAA..51..221Z", "2016AJ....151...38D", "2016IAUS..312..128D", "2016MNRAS.458.2694B", "2018ApJ...852...87D", "2018JOSS....3..899D", "2018PASA...35....5D", "2019JOSS....4..898D", "2019JPhCS1231a2031P", "2020Galax...8...60H", "2020MNRAS.493.2238A", "2021ApJS..257...58D", "2022ApJS..260...14D", "2022MNRAS.510.5444G", "2022MNRAS.511.1022D", "2024MNRAS.528..459P" ]
[ "astronomy" ]
8
[ "stars: Wolf-Rayet", "ISM: abundances", "planetary nebulae: individual: Abell 48", "Astrophysics - Solar and Stellar Astrophysics" ]
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[ "10.1093/mnras/stu203", "10.48550/arXiv.1403.0567" ]
1403
1403.0567_arXiv.txt
\label{a48:sec:introduction} The highly reddened planetary nebula Abell~48 (PN\,G029.0$+$00.4) and its central star (CS) have been the subject of recent spectroscopic studies \citep{Wachter2010,Depew2011,Todt2013,Frew2013}. The CS of Abell\,48 has been classified as Wolf--Rayet [WN5] \citep{Todt2013}, where the square brackets distinguish it from the massive WN stars. Abell 48 was first identified as a planetary nebula (PN) by \citet{Abell1955}. However, its nature remains a source of controversy whether it is a massive ring nebula or a PN as previously identified. Recently, \citet{Wachter2010} described it as a spectral type of WN6 with a surrounding ring nebula. But, \citet{Todt2013} concluded from spectral analysis of the CS and the surrounding nebula that Abell 48 is rather a PN with a low-mass CS than a massive (Pop I) WN star. Previously, \citet{Todt2010} also associated the CS of PB\,8 with [WN/C] class. Furthermore, IC\,4663 is another PN found to possess a [WN] star \citep{Miszalski2012}. A narrow-band H$\alpha$+[N\,\textsc{ii}] image of Abell~48 obtained by \citet{Jewitt1986} first showed its faint double-ring morphology. \citet{Zuckerman1986} identified it as a member of the elliptical morphological class. The H$\alpha$ image obtained from the SuperCOSMOS Sky H$\alpha$ Survey \citep{Parker2005} shows that the angular dimensions of the shell are about 46$\arcsec \times$ 38$\arcsec$, and are used throughout this paper. The first integral field spectroscopy of Abell~48 shows the same structure in the H$\alpha$ emission-line profile. But, a pair of bright point-symmetric regions is seen in [N\,\textsc{ii}] (see Fig.\,\ref{a48:fig3}), which could be because of the N$^{+}$ stratification layer produced by the photoionization process. A detailed study of the kinematic and ionization structure has not yet been carried out to date. This could be due to the absence of spatially resolved observations. \begin{table} \caption{Journal of the IFU observations with the ANU 2.3-m Telescope.} \label{a48:tab:obs:journal} \centering \begin{tabular}{lcccc} \hline \hline PN & Date ({\sc ut}) & $\lambda$ range ({\AA}) & $R$ & Exp.(s) \\ \hline Abell 48 & 2010/04/22 & 4415--5589 & 7000 & 1200\\ & & 5222--7070 & 7000 & 1200\\ & 2012/08/23 & 3295--5906 & 3000 & 1200\\ & & 5462--9326 & 3000 & 1200\\ \hline \end{tabular} \end{table} The main aim of this study is to investigate whether the [WN] model atmosphere from \citet{Todt2013} of a low-mass star can reproduce the ionization structure of a PN with the features like Abell~48. We present integral field unit (IFU) observations and a three-dimensional photoionization model of the ionized gas in Abell~48. The paper is organized as follows. Section \ref{a48:sec:observations} presents our new observational data. In Section~\ref{a48:sec:kinematic} we describe the morpho-kinematic structure, followed by an empirical analysis in Section~\ref{a48:sec:empirical}. We describe our photoionization model and the derived results in Sections~\ref{a48:sec:photoionization} and \ref{a48:sec:modelresults}, respectively. Our final conclusion is stated in Section\,\ref{a48:sec:conclusions}.
\label{a48:sec:conclusions} We have constructed a photoionization model for the nebula of Abell~48. This consists of a dense hollow cylinder, assuming homogeneous abundances. The three-dimensional density distribution was interpreted using the morpho-kinematic model determined from spatially resolved kinematic maps and the ISW model. Our aim was to construct a model that can reproduce the nebular emission-line spectra, temperatures and ionization structure determined from the observations. We have used the non-LTE model atmosphere from \citet{Todt2013} as the ionizing source. Using the empirical analysis methods, we have determined the temperatures and the elemental abundances from CELs and ORLs. We notice a discrepancy between temperatures estimated from $[$O~{\sc iii}$]$ CELs and those from the observed He\,{\sc i} ORLs. In particular, the abundance ratios derived from empirical analysis could also be susceptible to inaccurate values of electron temperature and density. However, we see that the predicted ionic abundances are in decent agreement with those deduced from the empirical analysis. The emission-line fluxes obtained from the model were in fair agreement with the observations. We notice large discrepancies between He\,{\sc i} electron temperatures derived from the model and the empirical analysis. The existence of clumps and low-ionization structures could solve the problems \citep{Liu2000}. Temperature fluctuations have been also proposed to be responsible for the discrepancies in temperatures determined from CELs and ORLs \citep{Peimbert1967,Peimbert1971}. Previously, we also saw large ORL--CEL abundance discrepancies in other PNe with hydrogen-deficient CSs, for example Abell~30 \citep{Ercolano2003b} and NGC~1501 \citep{Ercolano2004}. A fraction of H-deficient inclusions might produce those discrepancies, which could be ejected from the stellar surface during a very late thermal pulse (VLTP) phase or born-again event \citep[]{Iben1983}. However, the VLTP event is expected to produce a carbon-rich stellar surface abundance \citep{Herwig2001}, whereas in the case of Abell~48 negligible carbon was found at the stellar surface \citep[C/He~=~$3.5\times10^{-3}$ by mass;][]{Todt2013}. The stellar evolution of Abell~48 still remains unclear and needs to be investigated further. But, its extreme helium-rich atmosphere (85 per cent by mass) is more likely connected to a merging process of two white dwarfs as evidenced for R Cor Bor stars of similar chemical surface composition by observations \citep{Clayton2007,Garcia-Hernandez2009} and hydrodynamic simulations \citep{Staff2012,Zhang2012,Menon2013}. We derived a nebula ionized mass of $\sim0.8$~M$_{\bigodot}$. The high C/O ratio indicates that it is a predominantly C-rich nebula. The C/H ratio is largely over-abundant compared to the solar value of \citet{Asplund2009}, while oxygen, sulphur and argon are under-abundant. Moreover, nitrogen and neon are roughly similar to the solar values. Assuming a sub-solar metallicity progenitor, a 3rd dredge-up must have enriched carbon and nitrogen in AGB-phase. However, extremely high carbon must be produced through mixing processing in the He-rich layers during the He-shell flash. The low N/O ratio implies that the progenitor star never went through the hot bottom burning phase, which occurs in AGB stars with initial masses more than 5M$_{\bigodot}$ \citep{Karakas2007,Karakas2009}. Comparing the stellar parameters found by the model, $T_{\rm eff}$\,=\,70\,kK and $L_{\rm \star}/$L$_{\bigodot}$=~5500, with VLTP evolutionary tracks from \citet{Bloecker1995b}, we get a current mass of $\sim 0.62 {\rm M}_{\bigodot}$, which originated from a progenitor star with an initial mass of $\sim 3 {\rm M}_{\bigodot}$. However, the VLTP evolutionary tracks by \citet{MillerBertolami2006a} yield a current mass of $\sim 0.52 {\rm M}_{\bigodot}$ and a progenitor mass of $\sim1{\rm M}_{\bigodot}$, which is not consistent with the derived nebula ionized mass. Furthermore, time-scales for VLTP evolutionary track \citep{Bloecker1995b} imply that the CS has a post-AGB age of about $\sim$\,9\,000 yr, in agreement with the nebula's age determined from the kinematic analysis. We therefore conclude that Abell~48 originated from an $\sim3$~M$_{\bigodot}$ progenitor, which is consistent with the nebula's features.
14
3
1403.0567
Recent observations reveal that the central star of the planetary nebula Abell 48 exhibits spectral features similar to massive nitrogen-sequence Wolf-Rayet stars. This raises a pertinent question, whether it is still a planetary nebula or rather a ring nebula of a massive star. In this study, we have constructed a three-dimensional photoionization model of Abell 48, constrained by our new optical integral field spectroscopy. An analysis of the spatially resolved velocity distributions allowed us to constrain the geometry of Abell 48. We used the collisionally excited lines to obtain the nebular physical conditions and ionic abundances of nitrogen, oxygen, neon, sulphur and argon, relative to hydrogen. We also determined helium temperatures and ionic abundances of helium and carbon from the optical recombination lines. We obtained a good fit to the observations for most of the emission-line fluxes in our photoionization model. The ionic abundances deduced from our model are in decent agreement with those derived by the empirical analysis. However, we notice obvious discrepancies between helium temperatures derived from the model and the empirical analysis, as overestimated by our model. This could be due to the presence of a small fraction of cold metal-rich structures, which were not included in our model. It is found that the observed nebular line fluxes were best reproduced by using a hydrogen-deficient expanding model atmosphere as the ionizing source with an effective temperature of T<SUB>eff</SUB> = 70 kK and a stellar luminosity of L<SUB>⋆</SUB> = 5500 L<SUB>⊙</SUB>, which corresponds to a relatively low-mass progenitor star (∼3 M<SUB>⊙</SUB>) rather than a massive Pop I star.
false
[ "Abell", "helium temperatures", "spectral features", "our new optical integral field spectroscopy", "hydrogen", "massive nitrogen-sequence Wolf-Rayet stars", "a massive star", "a hydrogen-deficient expanding model atmosphere", "our photoionization model", "nitrogen", "ionic abundances", "decent agreement", "our model", "the model", "the central star", "the optical recombination lines", "a massive Pop", "kK", "the observed nebular line fluxes", "the empirical analysis" ]
9.983027
10.18394
151
745670
[ "Shimwell, Timothy W.", "Brown, Shea", "Feain, Ilana J.", "Feretti, Luigina", "Gaensler, B. M.", "Lage, Craig" ]
2014MNRAS.440.2901S
[ "Deep radio observations of the radio halo of the bullet cluster 1E 0657-55.8" ]
78
[ "CSIRO Astronomy &amp; Space Science, Australia Telescope National Facility, PO Box 76, Epping, NSW 1710, Australia", "Department of Physics and Astronomy, University of Iowa, 203 Van Allen Hall, Iowa City, IA 52242, USA", "CSIRO Astronomy &amp; Space Science, Australia Telescope National Facility, PO Box 76, Epping, NSW 1710, Australia", "INAF - Istituto di Radioastronomia, via Gobetti 101, I-40129 Bologna, Italy", "Sydney Institute for Astronomy, School of Physics, The University of Sydney, NSW 2006, Australia", "Center for Cosmology and Particle Physics, Department of Physics, New York University, New York, NY 10003, USA" ]
[ "2014PhyU...57..317V", "2015A&A...582A..20B", "2015ApJ...812..153O", "2015ApJ...813...77Y", "2015ApJ...815..116F", "2015MNRAS.449.1486S", "2015MNRAS.452.1328M", "2015arXiv150505790C", "2016A&A...595A.116M", "2016Ap&SS.361..255M", "2016ApJ...818..204V", "2016ApJ...820L..20D", "2016MNRAS.455.2402S", "2016MNRAS.458.2584B", "2016MNRAS.459..277S", "2016PASJ...68S..20U", "2017ApJ...845...81P", "2017MNRAS.464.2896C", "2017MNRAS.469.1476S", "2017MNRAS.471.1107H", "2017MNRAS.471.4747C", "2017MNRAS.472.2042P", "2017NatSR...716918M", "2018ApJ...852...65R", "2018ApJ...856..162W", "2018ApJ...857...79J", "2018ApJ...865...24D", "2018MNRAS.475.2743H", "2018MNRAS.475.2870C", "2018MNRAS.476.3415W", "2018MNRAS.476.5591B", "2018MNRAS.479..553S", "2018MNRAS.479..730Z", "2018arXiv180402588P", "2019A&A...622A..20H", "2019A&A...622A..21H", "2019ApJ...873...64D", "2019ApJ...876..154N", "2019ApJ...885...80N", "2019ApJ...887...26O", "2019MNRAS.487.4775B", "2019SSRv..215...16V", "2020A&A...634A..64B", "2020A&A...636A..30R", "2020A&A...642L...3Z", "2020ApJ...897...93B", "2020PASA...37...40W", "2020arXiv200601833W", "2021A&A...646A.135R", "2021A&A...654A..41R", "2021A&A...656A.154H", "2021ApJ...914...73Z", "2021ApJS..257...42L", "2021NatSR..1120087C", "2021PASA...38...10D", "2021PASA...38...31D", "2022A&A...657A..56K", "2022A&A...659A.146D", "2022A&A...663A.105P", "2022A&A...668A.107D", "2022AJ....164...83C", "2022ApJ...934...49G", "2022MNRAS.509.5821R", "2022MNRAS.511.3389V", "2022MNRAS.514.4969V", "2022MNRAS.515.1871R", "2022MNRAS.517.1299Z", "2022MNRAS.517.5232R", "2023A&A...669A...1R", "2023A&A...674A..53B", "2023ApJ...942...79T", "2023MNRAS.518.4595S", "2024A&A...686A...5B", "2024MNRAS.527.1194K", "2024arXiv240204876W", "2024arXiv240403944S", "2024arXiv240509384O", "2024arXiv240615750M" ]
[ "astronomy" ]
6
[ "acceleration of particles", "radiation mechanisms: non-thermal", "shock waves", "galaxies: clusters: individual: 1E 0657-55.8", "galaxies: clusters: intracluster medium", "radio continuum general", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
[ "1980ApJ...239L..93D", "1993ApJ...406..399G", "1995ASPC...77..433S", "1998A&A...332..395E", "2000ApJ...544..686L", "2001A&A...369..441G", "2001A&A...373..106F", "2001ApJ...557..560P", "2001MNRAS.320..365B", "2001MNRAS.328L..21L", "2004rcfg.proc..335K", "2005A&A...430L...5G", "2005A&A...440..867G", "2005A&A...441.1217B", "2006ApJ...648L.109C", "2006ESASP.604..723M", "2007APh....28..232K", "2008A&A...486..347G", "2008SSRv..134...93F", "2009A&A...503..409H", "2009A&A...503..707B", "2009A&A...507.1257G", "2009ApJ...691..963C", "2009ApJ...704.1349O", "2010MNRAS.405..155O", "2010Sci...330..347V", "2011A&A...525A.104P", "2011A&A...532A..71R", "2011ApJ...728...82M", "2011ApJ...733...63R", "2011MNRAS.412....2B", "2011MNRAS.415.2699A", "2011MNRAS.416..832W", "2012A&A...542A..93O", "2012A&ARv..20...54F", "2012ApJ...756...97K", "2012MNRAS.421.3300O", "2012SSRv..166..187B", "2012mgm..conf..397M", "2013A&A...551A..24V", "2013A&A...554A.102G", "2013ApJ...764...82B", "2013MNRAS.429.3564D", "2013MNRAS.433..812O", "2013MNRAS.433.1701O", "2013MNRAS.433.2920A", "2013PASJ...65...16A" ]
[ "10.1093/mnras/stu467", "10.48550/arXiv.1403.2393" ]
1403
1403.2393.txt
Low surface brightness, steep spectrum, diffuse radio emission has been observed to be associated with the intra-cluster medium (ICM) of many merging galaxy cluster systems (see \citealt{Bruggen_2012}, \citealt{Feretti_2012} and \citealt{Ferrari_2008} for recent reviews). The emission is due to synchrotron radiation from relativistic electrons gyrating in a magnetic field. When coincident with the centre of the merging cluster system, this emission is referred to as a radio halo, whereas if located on the periphery of the cluster it is known as a radio relic -- \cite{Kempner_2004} provides details on the classification of extended radio sources in galaxy clusters. While radio halos are believed to be the result of cluster-wide processes such as post-merger turbulence (see e.g. \citealt{Brunetti_2001} and \citealt{Petrosian_2001}) or proton-proton collisions (see e.g. \citealt{Dennison_1980}), radio relics are believed to be associated in some way with localised, post-merger shock-fronts (\citealt{Ensslin_1998}). Measurements of X-ray shocks in the regions of radio relics (see e.g. \citealt{Macario_2011}, \citealt{Bourdin_2013}, \citealt{Giacintucci_2008}, \citealt{Ogrean_2013a}, \citealt{Akamatsu_2013} and \citealt{Ogrean_2013b}) and synchrotron radio spectral index gradients (see e.g. \citealt{Weeren_2010}) have provided strong support to the theory that these objects are formed by first-order Fermi acceleration (diffusive shock acceleration). The exact connection/relationship between relics and halos is unknown, but some observations of radio halos have shown that weak X-ray shocks are approximately cospatial with edges of radio halos (see e.g. \citealt{Brown_2011} and \citealt{Markevitch_2012}). An example of such a system is 1E 0657-55.8 (usually referred to as `the bullet cluster' because of its distinctive morphology). In this case the radio images of the powerful radio halo indicate that it has an irregular shape (\citealt{Liang_2000}) and that the position of the radio edge and the X-ray shock is similar, however, it was unclear from the halo observations by \cite{Liang_2000} whether there was any correspondence between the morphology of the distinctive X-ray bow shock and the (indistinct) radio halo edge. We have obtained observations that are deeper than those of \cite{Liang_2000} and are able to characterise the halo in greater detail and at higher resolution. The primary aim of this study is to determine whether shocks influence the structure of radio halos. The bullet cluster with its prominent bow shock and bright radio halo is an obvious source for investigating the relationship between shocks and radio halos. We have used the Australia Telescope Compact Array (ATCA) to obtain deep, large fractional bandwidth (1.1-3.1\,GHz), polarimetric observations of the bullet cluster which we use to compare the radio emission from relativistic electrons with the X-ray emission from thermal gas. Hereafter we assume a concordance $\rm{\Lambda}$CDM cosmology, with $\rm{\Omega_{m}}$ = 0.3, $\rm{\Omega_\Lambda}$ = 0.7 and h $\equiv$ H$_{0}$/(100 km\,s$^{-1}$Mpc$^{-1}$) = 0.7. At $z=0.296$, the luminosity distance of 1E\,0657-55.8 is 1529 Mpc and 1$\arcsec$ corresponds to 4.413\,kpc. All coordinates are given in J2000.
We have analysed deep, wide-band, polarimetric, radio observations of the bullet cluster. After removing the contamination from radio sources, we have presented images of the bullet cluster radio halo that have significantly better sensitivity than previously published images of this halo. From these observations we have been able to characterise the morphology and spectral properties of the radio halo, but have been unable to detect polarised emission from any regions of the halo. From our study we have found the following: % \begin{itemize} \item We determined an integrated 1.3\,GHz flux-density of 52.5$\pm$2.1\,mJy from the region of significant radio halo emission. In this region our measurements are in good agreement with the \cite{Liang_2000} measurements from a comparable region. However, when integrating over a region much larger than the area of significant emission we measure an integrated 1.3\,GHz flux density of 56.4$\pm$2.3\,mJy which is lower than the large area \cite{Liang_2000} measurement of 78$\pm$5\,mJy. \item Our measured 1.3\,GHz to 3.1\,GHz spectral index ($-1.57\pm0.05$) is steeper than the 0.8\,GHz to 8.8\,GHz spectral index measured by \cite{Liang_2000} (-1.3$\pm$0.1). \item The radio emission extends as far as the observed X-ray emission in the direction of the cluster merger, but perpendicular to this axis, the extent of the radio emission is smaller than that of the X-ray emission. \item We find evidence for two peaks in the radio halo emission. One is coincident with the X-ray centroid of the main cluster and another is close to the bullet. \item The centroids of surface brightness slices through the halo and X-ray structures are better aligned in the direction perpendicular to the merging axis than in the direction of the cluster merger. \item There is a distinctive edge to the bullet cluster radio halo, which is coincident with the bow shock detected in deep X-ray observations. \item The radio intensity and spectral index distribution do not have tight correspondences with the X-ray brightness, X-ray temperature or weak-lensing mass reconstruction. \item The radio spectral index varies across the cluster but there is no clear trend. \item We do not detect any polarised emission from the radio halo. Our 5$\sigma_{Q,U}$ upper limits on the fractional polarisation in the centre of the cluster over the 1.1-3.1\,GHz band are 13\% at a resolution of 21$\arcsec$. \end{itemize}
14
3
1403.2393
We present deep 1.1-3.1 GHz Australia Telescope Compact Array observations of the radio halo of the bullet cluster, 1E 0657-55.8. In comparison to existing images of this radio halo, the detection in our images is at higher significance. The radio halo is as extended as the X-ray emission in the direction of cluster merger but is significantly less extended than the X-ray emission in the perpendicular direction. At low significance, we detect a faint second peak in the radio halo close to the X-ray centroid of the smaller sub-cluster (the bullet) suggesting that, similarly to the X-ray emission, the radio halo may consist of two components. Finally, we find that the distinctive shape of the western edge of the radio halo traces out the X-ray detected bow shock. The radio halo morphology and the lack of strong point-to-point correlations between radio, X-ray and weak-lensing properties suggest that the radio halo is still being formed. The colocation of the X-ray shock with a distinctive radio brightness edge illustrates that the shock is influencing the structure of the radio halo. These observations support the theory that shocks and turbulence influence the formation and evolution of radio halo synchrotron emission.
false
[ "radio halo synchrotron emission", "-", "the X-ray emission", "the X-ray shock", "ray", "the X-ray centroid", "cluster merger", "The radio halo morphology", "bow shock", "The radio halo", "the radio halo", "this radio halo", "cluster", "a distinctive radio brightness edge", "higher significance", "radio, X-ray and weak-lensing properties", "point", "low significance", "existing images", "second" ]
14.391587
4.751124
152
788622
[ "Singaravelu, Balasubramanian", "Cabanac, Remi A." ]
2014PASP..126..386S
[ "Obstructed Telescopes Versus Unobstructed Telescopes for Wide Field Survey—A Quantitative Analysis" ]
8
[ "-", "-" ]
[ "2014SPIE.9143E..4XH", "2017AJ....154..244T", "2017ApOpt..56.8639M", "2019MNRAS.488.5057L", "2020OExpr..2837958L", "2021ApOpt..60.6438L", "2021ExA....51.1623B", "2023OExpr..3134477L" ]
[ "astronomy" ]
9
[ "Astrophysics - Instrumentation and Methods for Astrophysics" ]
[ "1958JOSA...48..844T", "1965poet.book.....B", "1972ApOpt..11.2986K", "1978ApJ...222....1K", "1978ApOpt..17.2677R", "1980ApOpt..19.3640K", "1987sdap.book.....S", "1996A&AS..117..393B", "1998A&AS..127....1L", "2005ESASP.576...15P", "2009A&A...504..359J", "2009MmSAI..80..422B", "2010AJ....139.2097P", "2010MNRAS.404..350R", "2010Natur.464.1018S", "2010PASP..122..439R", "2010SPIE.7731E..1GL", "2012ExA....34..311T", "2012MNRAS.421.1385V" ]
[ "10.1086/676317", "10.48550/arXiv.1403.2879" ]
1403
1403.2879_arXiv.txt
Many astrophysical observations demand very high-resolution imaging and a stable PSF over the desired field--of--view. The former is limited by the size of the pupil and presence of central obstruction, while the latter is limited only by the obstruction. Theoretically unobstructed pupils deliver a simpler PSF compared to obstructed designs. We know from the studies of \citet*{row10} that the accurate knowledge of the PSF is significant to make precise science measurements on astronomical images. Recent works done by \citet*{lam10} and \citet*{lev11} suggest that unobstructed pupil have a faster survey rate for wide-field survey missions to study dark energy. These wide-survey images are used to perform strong lensing and weak lensing analyses. We know that lensing analyses depends on the precision with which we can measure the ellipticity and position--angle of a large number background sources and hence the shear and convergence produced by the lensing mass. The objective of this paper is to quantify the image quality of an obstructed and unobstructed telescope for such a wide-field survey mission. In this paper we focus on space based telescopes because ground based telescopes even with `multi conjugate adaptive optics' (MCAO) are still not suitable to perform high resolution surveys. \subsection{Three Mirror Anastigmat Telescopes} Optical designs based on three aspheric mirrors (dubbed Three--Mirror--Anastigmat or TMA) are favored in modern astronomical instrumentation because these telescopes have a wider field--of--view for a given pupil size. And it is possible to build a range of diffraction-limited designs using this configuration. TMA telescopes were proposed by \citet*{pau35} and developed by \citet*{bak69} and \citet*{kor72}. Traditionally, these telescopes consist of a secondary mirror which partially obstructs the light falling on the primary mirror. This type of three mirror telescope are called obstructed TMA (hereafter OTMA). A few telescopes have been built recently in which the secondary mirror is offset from the path of the incoming light and no obstruction is present in front of the primary reflector. These type of telescopes are called unobstructed TMA (hereafter UTMA). The design of UTMA telescopes was discussed by \citet*{coo79} and \citet*{kor80}. For stellar astrometry, UTMA design delivers better precision, for example GAIA \citep*{per05}. Also unobstructed telescopes can image exoplanets close to bright stars with high contrast \citep*{ser10}. Hence, UTMA designs have been proposed for exoplanet characterization missions like EChO \citep*{tin12} and SPICES \citep*{boc12}. In the next section we briefly summarize the theoretical advantage of using an unobstructed pupil. \subsection{Advantage of an unobstructed pupil} The effective light gathering area of an obstructed pupil is lesser than an unobstructed pupil. And the blur size on the image plane increases with increasing obstruction. The focal image of a point source at infinity or PSF of a focusing optical system is the Fourier transform of its pupil shape. In the case of a circular pupil, the shape of the PSF is the Airy pattern with a bright central disc and faint concentric rings. We know from the studies of \citet*{tay58} that in the Airy pattern of an annular aperture there is significant transfer of energy from the central disc to the outer rings. The presence of support structure will also produce additional artefacts on the PSF. \citet*{lev11} show in their work that the survey rate for a mission is directly proportional to the light gathering area and inversely proportional to the blur size for a given SNR. Therefore unobstructed telescopes must have faster survey speed for a given SNR. They also demonstrate that for unobstructed telescopes there is an increased density of resolved galaxies for weak lensing analyses. \citet*{lam10} show that the diffraction pattern of an OTMA telescope can destroy or mimic the lensing signals we desire to study because of the larger blur size. Hence the tighter PSF of UTMA telescopes should be beneficial for wide-field lensing missions. In this work we will quantify the precision with which UTMA telescopes can perform science measurements for weak lensing analyses compared to OTMA telescopes having same characteristics. \subsection{Outline} The objective of this work is to quantify the precision with which UTMA and OTMA telescopes can perform morphological parameter measurements for a wide field survey. In order to reach this goal the work was organized as follows. First, we modeled and optimized in parallel an OTMA and an UTMA telescope, both having the same primary mirror, effective focal length and FoV. The science requirements for our study are same as that of the {\it Euclid} mission \citep*{euc11}. Second, we created an end-to-end semi-realistic samples of elliptical galaxies passed through the full instrumental path, PSF convolution, CCD pixelisation, and noise effects. Third we selected and controlled the biases of the fitting routine for measuring galaxy morphologies. Then we measured the ellipticity in the simulated galaxies and calculated the error introduced by the PSF on the error budget. Finally, we compare the precision with which both the designs can perform the desired measurements. The paper structure sequentially follows these steps.
The significant results of our analysis are as follows. In case study \#1 and \#3 we show that the intrinsic properties of the OTMA PSF, like the bright concentric rings and diffraction spikes did not affect the morphological parameter measurements significantly, if the PSF is reconstructed within 10 arc-minutes of the source. If the PSF is reconstructed beyond 10 arc-minute the performance of the OTMA design degrades. But case study \#2 shows that the precision of the UTMA design is not affected even if the PSF is reconstructed 50 arc-minutes from the point of interest. This is because the PSF of UTMA design is homogeneous over the desired FoV compared to the OTMA design.
14
3
1403.2879
Telescopes with unobstructed pupil are known to deliver clean point spread function (PSF) to their focal plane, in contrast to traditional telescopes with obstructed pupil. Recent progress in the manufacturing aspheric surfaces and mounting accuracy favors unobstructed telescopes over obstructed telescopes for science cases that demand stable and clean PSF over the entire field-of-view. In this paper we compare the image quality of an unobstructed Three-Mirror-Anastigmat (TMA) design with that of an obstructed TMA. Both the designs have the same primary mirror, effective focal length, field-of-view and detector characteristics. We demonstrate using simulated images of faint elliptical galaxies imaged through the two designs, that both the designs can measure morphological parameters with same precision, if the PSF is reconstructed within 12 arc-minutes of the source. We also demonstrate that, the unobstructed design delivers desirable precision even if the PSF is reconstructed 50 arc-minutes away from the source. Therefore the PSF of unobstructed design is uniform over a wider field-of-view compared to an obstructed design. The image quality is given by the 1$\sigma$ error-bars (68% confidence level) in the fitted values of the axis-ratio and position-angle of the simulated galaxies.
false
[ "unobstructed design", "obstructed telescopes", "obstructed pupil", "unobstructed telescopes", "PSF", "clean point spread function", "traditional telescopes", "same precision", "view", "unobstructed pupil", "effective focal length", "faint elliptical galaxies", "Telescopes", "desirable precision", "an obstructed design", "science cases", "TMA", "morphological parameters", "stable and clean PSF", "minutes" ]
14.180011
10.347031
10
810650
[ "Yasutake, Nobutoshi", "Łastowiecki, Rafał", "Benić, Sanjin", "Blaschke, David", "Maruyama, Toshiki", "Tatsumi, Toshitaka" ]
2014PhRvC..89f5803Y
[ "Finite-size effects at the hadron-quark transition and heavy hybrid stars" ]
96
[ "Department of Physics, Chiba Institute of Technology, 2-1-1 Shibazono, Narashino, Chiba 275-0023, Japan", "Institute for Theoretical Physics, University of Wrocław, Max Born pl. 9, 50-204 Wrocław, Poland", "Physics Department, Faculty of Science, University of Zagreb, Zagreb 10000, Croatia", "Institute for Theoretical Physics, University of Wrocław, Max Born pl. 9, 50-204 Wrocław, Poland; Bogoliubov Laboratory for Theoretical Physics, JINR Dubna, 141980 Dubna, Russia", "Advanced Science Research Center, Japan Atomic Energy Agency, Tokai, Ibaraki 319-1195, Japan", "Department of Physics, Kyoto University, Kyoto 606-8502, Japan" ]
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[ "astronomy", "physics" ]
8
[ "26.60.-c", "97.60.Jd", "21.65.-f", "12.39.-x", "Nuclear matter aspects of neutron stars", "Neutron stars", "Nuclear matter", "Phenomenological quark models", "Astrophysics - High Energy Astrophysical Phenomena", "High Energy Physics - Phenomenology", "Nuclear Theory" ]
[ "1971ApJ...170..299B", "1987MPLA....2..805H", "1987PhR...149....1M", "1990PhRvC..42.1965B", "1993PhRvL..70.1355H", "1995PhLB..355...21S", "1997PhRvC..56.1720P", "1998PhRvC..58.1804A", "1998PhRvC..58.3688B", "2000ApJ...533..406B", "2000PhRvC..61b5802V", "2000PhRvL..85.2048P", "2001A&A...368..561B", "2001hep.ph....4183G", "2002PhLB..541...93V", "2003A&A...408..675B", "2003NuPhA.718..359T", "2003NuPhA.723..291V", "2003PhRvC..68a5203B", "2003astro.ph.12446B", "2004A&A...416..991A", "2004NuPhA.744..227V", "2005MNRAS.361.1243P", "2005PThPh.113..953Y", "2005PhRvC..71d5801G", "2005PhRvC..72a5802M", "2006A&A...450..345K", "2006ApJ...639..382L", "2006PThPh.115..337E", "2006PhRvC..74b5803P", "2006PhRvD..73e4504P", "2007A&A...465..533Z", "2007PhLB..654..170K", "2007PhRvC..75f5804B", "2007PhRvD..75h4012Y", "2007PhRvD..76i4501K", "2007PhRvD..76l3015M", "2007PhRvL..99b2001I", "2007PrPNP..59..139B", "2008PhRvD..77f3004P", "2008PhRvD..77j3006N", "2008PhRvD..78k4002N", "2009MNRAS.392...52A", "2009PhRvC..80f5807B", "2009PhRvD..80g4029F", "2009PhRvL.102h1101S", "2010ApJ...722...33S", "2010JPhG...37i4063B", "2010Natur.467.1081D", "2010PThPS.186...63T", "2010PhRvD..82e4026C", "2010PhRvD..82l5018P", "2012PhRvC..86b5203P", "2012arXiv1203.5341R", "2012arXiv1208.0427Y", "2013A&A...551A..13C", "2013A&A...551A..61Z", "2013ApJ...762...96B", "2013ApJ...765L...5S", "2013PhRvC..88d5803L", "2013PhRvD..87b3001O", "2013PhRvD..87c6004M", "2013PhRvD..88h3013A", "2013PhRvD..88h5001K", "2013Sci...340..448A", "2014EPJA...50..111B", "2014JPhCS.496a2002B", "2014PhLB..736..241S", "2014PhRvD..89a6007B" ]
[ "10.1103/PhysRevC.89.065803", "10.48550/arXiv.1403.7492" ]
1403
1403.7492_arXiv.txt
Introduction} The equation of state~(EoS) is the central quantity for the study of compact stars. Since modern lattice QCD simulations are not applicable at large baryon densities and low temperatures $T\simeq 0$, there is a large uncertainty in theoretical descriptions of the behavior of matter at extreme densities. The understanding may be improved by studying astrophysical phenomena; namely, we may use the known astrophysical constraints from observations of compact stars in order to provide constraints on the EoS. Recently, the idea has been pursued to use a Bayesian analysis (BA) for ``inversion'' of the Tolman-Oppenheimer-Volkoff equations, i.e. to extract a probability measure for models of the cold EoS in the pressure-energy density plane from observational data related to masses and radii of compact stars. While first analyses of this type have favored burst sources with rather uncertain and model dependent statements about radii \cite{Steiner:2010fz,Steiner:2012xt}, a very recent BA uses a set of stronger and statistically independent observations, testing also the possibility of a first order phase transition at supersaturation densities \cite{Blaschke:2014via}. At this point the strongest restriction to the EoS is provided by the recent measurement of the high mass of $\sim 2~M_\odot$ from observations of the pulsars PSR J1614-2230 by Demorest et.~al.~\cite{Demorest:2010bx} and PSR J0348+0432 by Antoniadis et.~al.~\cite{Antoniadis:2013pzd}. The recent BA \cite{Blaschke:2014via} makes use of this constraint together with a new mass-radius constraint from the precise timing analysis of the nearest known millisecond pulsar PSR J0437-4715 \cite{Bogdanov:2012md} and the constraint on the gravitational binding for the neutron star B in the binary system J0737-3039(B) \cite{Podsiadlowski:2005ig}, see also \cite{Kitaura:2006}, at the precisely measured gravitational mass of $1.249 \pm 0.001~M_\odot$. There are many studies relating astrophysical phenomena involving compact stars and the properties of matter at extreme densities, eventually including the possibility of a quark deconfinement transition. These concern, e.~g., the cooling of compact stars~\cite{Page:2000wt,Blaschke:1999qx,Blaschke:2000dy,Grigorian:2004jq}, gravitational wave emission \cite{Lin:2005zda,Yasutake:2007st,Abdikamalov:2008df}, neutrino emission~\cite{Hatsuda:1987ck,Nakazato:2008su,Sagert:2008ka}, eigenfrequencies~\cite{Burgio:2003mr}, and the energy release during the collapse of neutron stars to quark stars~\cite{Aguilera:2002dh,Yasutake:2004kx,Zdunik:2006uw}. The study of the baryon-baryon (BB) interaction in lattice QCD simulations recently became a hot topic ~\cite{Ishii:2006ec,Walker-Loud:2014iea}. Experiments like JPARC will also provide valuable information on the BB interaction. In the near future the EoS in the hadronic phase may be determined by incorporating this information on the BB interaction in the Br\"{u}ckner-Hartree-Fock~(BHF) theory~\cite{bhf}, the variational approach~\cite{Akmal:1998cf,Takano:2010zz}, or the Dirac-Br\"{u}ckner-Hartree-Fock~(DBHF) theory~\cite{Brockmann:1990cn,vanDalen:2004pn}. In this paper we adopt the BHF theory for hadronic matter. Out of a several of possible models for quark matter we use the two flavor covariant non-local Nambu--Jona-Lasinio (nlNJL) model \cite{Contrera:2010kz,Benic:2013eqa} with vector interactions \cite{Alvarez-Castillo:2013spa}. The advantage over the usual local version of the NJL model is due to the introduction of the additional gradient self-energy channel and due to the explicit momentum dependence of all the dressing functions of the quark propagator. Both of these improvements are well founded on lattice QCD data \cite{Parappilly:2005ei,Kamleh:2007ud} and Dyson-Schwinger equation studies \cite{Bhagwat:2003vw, Fischer:2009gk, Roberts:2012sv}, and make the non-local NJL model a well-calibrated, effective low-energy QCD approach to the thermodynamics of quark matter. The main purpose of this work is to examine the features and the astrophysical consequences of the mixed phase between the pure quark and hadron matter phases by considering {\it finite-size effects}. Taking into account the surface tension and the charge screening we find the non-uniform, so-called ``pasta" structures at the hadron-quark interface. In this work we investigate more in detail the occurrence of pasta structures for the values of the surface tension $\sigma=10$ MeV fm$^{-2}$ and $40$ MeV fm$^{-2}$. For weak surface tension the EoS of the mixed phase becomes similar to the one of a bulk Gibbs construction, while for strong surface tension it approaches the result of a Maxwell construction \cite{Voskresensky:2001jq,Voskresensky:2002hu,Endo:2005zt,Maruyama:2007ey}, in which the maximum masses with the phase transition are around $1.5 M_\odot$ and a simple bag model was used for modeling the quark phase. This model gives a quite simple description of quark matter, and it should be replaced by a more sophisticated one to study more realistically the quark-hadron phase transition. This is the aim of the present work. We construct the hybrid EoS and the corresponding hybrid star sequences. For the calculation of the quark matter EoS we use the following values for the ratio of the vector and the scalar channel couplings $\eta_V = G_V/G_S = 0.10$ and $\eta_V=0.20$. Stable hybrid stars respecting the $2~M_\odot$ constraint are found in the case of $\eta_V=0.10$. The bulk Gibbs construction for this case supports only a mixed phase in the core. However, taking into account finite-size effects the cores of massive hybrid stars are composed of pure quark matter. For $\eta_V=0.20$ the $2M_\odot$ stars are mainly composed of hadron matter. With the appearance of quark matter at higher densities the star becomes unstable. For $\eta_V=0.10$ quark matter appears at low densities causing a reduction of the proton fraction at the onset of the mixed phase below the threshold value of $1/9$ for the onset of the direct URCA (dURCA) process in the $n-p-e$ phase, while for $\eta_V=0.20$ the proton fraction exceeds this value. This paper is organized as follows. In Sec.~\ref{msec:eos}, we outline our framework for obtaining the hybrid EoS with pasta phase. Sec.~\ref{sec:results} contains numerical results for the EoS with the different quark-hadron mixed phase constructions including the pasta phase as well as for the corresponding compact star sequences. Sec.~\ref{sec:conclude} is devoted to the conclusion and a discussion of some astrophysical implications of our results.
\label{sec:conclude} We have studied the hadron-to-quark-matter phase transition with {\it finite-size effects} by imposing the Gibbs conditions on the phase equilibrium, and calculated the density profiles in a self-consistent manner. For the quark phase we used the covariant nonlocal NJL model, while the hadron phase was given by the BHF EoS with Bonn-B potential. At strong surface tension, the EoS of the hadron-quark phase transition gets close to that given by the Maxwell construction. This is due to the mechanical instability of the geometrical structure induced by the surface tension. The pressure of the mixed phase shows a similar behavior to that of the bag model \cite{Voskresensky:2001jq,Voskresensky:2002hu,Endo:2005zt,Maruyama:2007ey}. It appears that this behavior of the hadron-quark phase transition is universal. Since the EoS has many uncertainties, especially concerning quark matter we plan to study this behavior using other quark and hadron models such as \cite{Chen:2013tfa}. Moreover, color superconductivity may also change our results \cite{Pagliara:2007ph,Klahn:2013kga}. We have found that the models used here describe compact star sequences with maximum masses exceeding the present constraint of $\sim 2M_\odot$ as deduced from observations \cite{Demorest:2010bx,Antoniadis:2013pzd}. For the low value of the vector coupling quarks appear at low densities, which might work in favor of suppressing the dURCA cooling channel in nuclear matter by the early onset of quark matter. If we conjecture that future observations would find quark matter in neutron stars our results would indicate that vector interactions in the quark phase are small, unlike the ones in nuclear matter. Since the phase transition to quark matter leads to a softening of the EoS, this is usually associated with the reduction in the maximum mass. The fact that the nuclear EoS employed in this work gives a neutron star surpassing $2M_\odot$ therefore works in favour of obtaining a sufficiently heavy hybrid star to exceeding $2M_\odot$ as well. However, we should keep in mind that this statement depends on the relative stiffness of the nuclear and the quark EoS at the highest densities reached in the core. There are notable exceptions in the literature, see e.~g.~\cite{Baldo:2003vx,Maruyama:2007ey,Lastowiecki:2011hh,Zdunik:2012dj,Alford:2013aca} where one finds an opposite scenario so that the maximum mass of the hybrid star is actually larger than the maximum mass of the pure nuclear star. In the context of the microscopically founded EoS model presented in this paper with the BHF approach to the hadronic phase and the nonlocal chiral quark model for the deconfined phase, the nonstrange hybrid star scenario appears as the most conservative one. Scenarios including strangeness in the hadronic and/or quark matter phase may require additional stiffening effects, that are beginning to be explored \cite{Benic:2014iaa}, in order to meet the $2M_\odot$ mass constraint. We shall return to such scenarios in subsequent work.
14
3
1403.7492
We study the role of finite-size effects at the hadron-quark phase transition in a new hybrid equation of state constructed from an ab initio Brückner-Hartree-Fock equation of state with the realistic Bonn-B potential for the hadronic phase and a covariant nonlocal Nambu-Jona-Lasinio model for the quark phase. We construct static hybrid star sequences and find that our model can support stable hybrid stars with an onset of quark matter below 2M<SUB>⊙</SUB> and a maximum mass above 2.17M<SUB>⊙</SUB> in agreement with recent observations. If the finite-size effects are taken into account the core is composed of pure quark matter. Provided that the quark vector channel interaction is small, and the finite size effects are taken into account, quark matter appears at densities 2-3 times the nuclear saturation density. In that case the proton fraction in the hadronic phase remains below the value required by the onset of the direct URCA process, so that the early onset of quark matter shall affect on the rapid cooling of the star.
false
[ "densities", "quark matter", "pure quark matter", "stable hybrid stars", "static hybrid star sequences", "the quark phase", "state", "the quark vector channel interaction", "recent observations", "2-3 times the nuclear saturation density", "a new hybrid equation", "the hadron-quark phase transition", "URCA", "the hadronic phase", "the early onset", "Bonn", "account", "finite-size effects", "the finite size effects", "the direct URCA process" ]
4.862336
1.830322
71
437455
[ "Török, Gabriel", "Urbanec, Martin", "Adámek, Karel", "Urbancová, Gabriela" ]
2014A&A...564L...5T
[ "Appearance of innermost stable circular orbits of accretion discs around rotating neutron stars" ]
9
[ "Institute of Physics, Faculty of Philosophy and Science, Silesian University in Opava, Bezručovo nám. 13, 74601, Opava, Czech Republic", "Institute of Physics, Faculty of Philosophy and Science, Silesian University in Opava, Bezručovo nám. 13, 74601, Opava, Czech Republic", "Institute of Physics, Faculty of Philosophy and Science, Silesian University in Opava, Bezručovo nám. 13, 74601, Opava, Czech Republic", "Institute of Physics, Faculty of Philosophy and Science, Silesian University in Opava, Bezručovo nám. 13, 74601, Opava, Czech Republic" ]
[ "2015GReGr..47...27S", "2016ApJ...833..273T", "2017LRR....20....7P", "2017PhLB..768..198D", "2017bhns.work...23G", "2018ApJ...861..141L", "2019ApJ...877...66U", "2019PhRvD..99b4041G", "2023PhRvD.108h4002C" ]
[ "astronomy" ]
4
[ "accretion", "accretion disks", "stars: neutron", "X-rays: binaries", "Astrophysics - High Energy Astrophysical Phenomena" ]
[ "1968ApJ...153..807H", "1972ApJ...178..347B", "1985Natur.316..239A", "1985Natur.317..681L", "1987Natur.327..303P", "1990ApJ...358..538K", "1998ApJ...508..791M", "1998PhRvD..58j4011S", "1999ApJ...520..763P", "1999PhRvL..82...17S", "2001A&A...374L..19A", "2001ApJ...550..426L", "2002A&A...381L..21A", "2002ApJ...577L..23T", "2003MNRAS.344..978R", "2003gr.qc....12070A", "2005A&A...439L..27P", "2005AJ....129.1993M", "2005ChJAA...5...41Z", "2005MNRAS.358..923B", "2005MNRAS.361..855B", "2006MNRAS.370.1140B", "2006PhRvD..73j4038P", "2007A&A...467..641S", "2007AcA....57....1A", "2007PhR...442..109L", "2008A&A...489..963S", "2008AcA....58....1T", "2008CQGra..25v5016K", "2008MNRAS.391.1472A", "2009CQGra..26e5011S", "2010A&A...522A..72U", "2010ATel.2919....1B", "2010AcA....60..149U", "2010ApJ...714..748T", "2010ApJ...717L.149M", "2010CQGra..27d5001B", "2012ARNPS..62..485L", "2012ApJ...756...82P", "2012CQGra..29f5012B", "2013MNRAS.433.1903U", "2013MNRAS.434.2825K", "2013MNRAS.435.3494W", "2013Sci...340..448A", "2021ASSL..461..143P" ]
[ "10.1051/0004-6361/201423541", "10.48550/arXiv.1403.3728" ]
1403
1403.3728_arXiv.txt
\label{section:introduction} In Newtonian theory, circular trajectories of test particles orbiting around a spherical central body of mass $M$ are stable to small radial perturbations at any external radii $r$. The same trajectories calculated using a general relativistic description exhibit an instability below the critical radius of the marginally stable cicular orbit $r_{\mathrm{ms}}$ \citep[e.g.,][]{bar-etal:1972}. This radius is frequently considered as the innermost stable circular orbit (ISCO) of an accretion disc that orbits a black hole or a neutron star (NS), $r_{\mathrm{ISCO}}= r_{\mathrm{ms}}$. The ISCO is often assumed as a unique prediction of Einstein's general relativity. Several decades after the ISCO concept has been introduced, it was argued that ISCO also appears around highly elliptic bodies described by the Newtonian theory \citep[see works of][]{zdu-gou:2001,ams-etal:2002,klu-ros:2013}. In this sense, several phenomena related to rotating oblate neutron stars can be well understood in terms of the interplay between the effects of general relativity and Newtonian theory. Namely, as in the Lense-Thirring or Kerr spacetime, the position of ISCO decreases with growing angular momentum $j$ of the neutron star (given by the NS spin). At the same time, it increases with the increasing influence of the NS quadrupole moment $q$ \citep[which determines the oblateness parameter $\tilde{q}\equiv q/j^2$, see, e.g.,][]{urb-etal:2013}. In this paper we explore the consequences of this interplay on the behaviour of NS compactness. Finally, we focus on low-mass X-ray binary systems (LMXBs) and find astrophysical implications for the distribution of neutron stars with an ISCO (ISCO-NS). \begin{figure*}[t] \begin{center} \includegraphics[width=1\hsize]{f1.eps} \end{center} \caption{\small{Left: Dependence of the geodesic ISCO radius $r_{\mathrm{ISCO}}= r_{\mathrm{ms}}$ on the NS angular momentum $j$ and the oblateness factor $\tilde{q}$. The shaded area indicates the Kerr-like region where $r_{\mathrm{ISCO}}$ decreases with increasing $j$. We note that for the concrete NS models we assume here, a fixed value of $\tilde{q}$ corresponds to a fixed NS central density $\rho$. Moreover, in the limit of $\tilde{q}=1$ the NS spacetime approaches the Kerr spacetime and values of $\tilde{q}\in(1,\,3)$ correspond to the highest possible values of $\rho$ (the oblateness parameter $\tilde{q}$ is also called the Kerr parameter). Right: Dependence of compactness $K$ on the NS spin frequency $\nu_{\mathrm{spin}}$ and mass $M$. The shaded area indicates the region of $K<1$ where the ISCO does not appear above the NS surface.}} \label{figure:risco} \end{figure*}
\label{section:conclusions} Our analysis clearly revealed the non-monotonicity of the dependence of the compactness factor $K$ on NS spin. The occurence of NS spin interval in which the ISCO appears for two very different ranges of NS mass (and the inferred double-peaked ISCO-NS spin distribution) depends on the detailed behaviour of $K$. For the four EoS discussed here, this behaviour is rather similar. However, normalization of $K$ and its exact dependence on $\nu_{\mathrm{spin}}$ and $M$ depend on the particular EoS. More detailed investigation that takes into account a large set of EoS needs to be performed to make better astrophysical assessments. Moreover, we assumed here the simplified identity $r_{\mathrm{ISCO}}= r_{\mathrm{ms}}$, but various effects might lead to a decrease or increase of $K$ (e.g. viscosity, pressure forces, or magnetic field influence), and the chosen spacetime description can play some role as well \citep[e.g.][]{alp-psa:2008,str-sra:2009,bak-etal:2010,bak-etal:2012,kot-etal:2008}. Despite these uncertainties, which need to be adressed in the subsequent analysis, our findings can be useful in several astrophysical applications. {Clearly, the predictions on the NS parameters would be strong if the ISCO presence or absence in the individual systems were observationally confirmed. Moreover, one can speculate on some consequences for evolution of the spinning-up sources.} {Although more development is required, some implications can be consired immediately. One possibility is to use our results in a study aimed to distinguish between different models of high-frequency quasiperiodic oscillations (HF QPOs). There is no consensus as yet on the HF QPO origin, and various models have been proposed \citep[e.g.,][and several others]{alp-sha:1985,lam-etal:1985,mil-etal:1998,psa-etal:1999,ste-vie:1999,abr-klu:2001,tit-ken:2002,rez-etal:2003,pet:2005,zha:2005,sra-etal:2007,stu-etal:2008}. Among numerous ideas,} it was suggested that the strong luminosity variations observed on the kHz time-scales in the neutron star X-ray binaries result from the modulation of accretion flow that heats the boundary layer between the flow and the NS surface \citep[][]{pac:1987,hor:2005,abr-etal:2007}. Motivated by this suggestion, \cite{urb-etal:2010} assumed the Paczynski modulation mechanism for the specific epicyclic resonance QPO model and investigated the implied restrictions on the NS mass and angular momentum. Taking into account our present findings and the general idea that (any) disc oscillations are responsible for exciting the Paczynski modulation mechanism, one might expect that the spin distribution of the HF QPO sources is double-peaked when NS masses are distributed around the intermediate value $M_0\sim1.4M_{\sun}$.
14
3
1403.3728
The innermost stable cicular orbit (ISCO) of an accretion disc that orbits a neutron star (NS) is often assumed to be a unique prediction of general relativity. However, it has been argued that ISCO also appears around highly elliptic bodies described by Newtonian theory. In this sense, the behaviour of an ISCO around a rotating oblate neutron star is formed by the interplay between relativistic and Newtonian effects. Here we briefly explore the consequences of this interplay using a straightforward analytic approach as well as numerical models that involve modern NS equations of state. We examine the ratio K between the ISCO radius and the radius of the neutron star. We find that, with growing NS spin, the ratio K first decreases, but then starts to increase. This non-monotonic behaviour of K can give rise to a neutron star spin interval in which ISCO appears for two very different ranges of NS mass. This may strongly affect the distribution of neutron stars that have an ISCO (ISCO-NS). When (all) neutron stars are distributed around a high mass M<SUB>0</SUB>, the ISCO-NS spin distribution is roughly the same as the spin distribution corresponding to all neutron stars. In contrast, if M<SUB>0</SUB> is low, the ISCO-NS distribution can only have a peak around a high value of spin. Finally, an intermediate value of M<SUB>0</SUB> can imply an ISCO-NS distribution divided into two distinct groups of slow and fast rotators. Our findings have immediate astrophysical applications. They can be used for example to distinguish between different models of high-frequency quasiperiodic oscillations observed in low-mass NS X-ray binaries.
false
[ "NS mass", "growing NS spin", "NS X-ray", "modern NS equations", "NS", "neutron stars", "ISCO", "low-mass NS X-ray binaries", "M<SUB>0</SUB", "Newtonian theory", "spin", "general relativity", "ISCO-NS", "a neutron star spin interval", "the ISCO-NS spin distribution", "different models", "a high mass M<SUB>0</SUB", "numerical models", "Newtonian", "an ISCO-NS distribution" ]
6.69564
2.581272
62
482814
[ "Lis, D. C.", "Schilke, P.", "Bergin, E. A.", "Gerin, M.", "Black, J. H.", "Comito, C.", "De Luca, M.", "Godard, B.", "Higgins, R.", "Le Petit, F.", "Pearson, J. C.", "Pellegrini, E. W.", "Phillips, T. G.", "Yu, S." ]
2014ApJ...785..135L
[ "Widespread Rotationally Hot Hydronium Ion in the Galactic Interstellar Medium" ]
20
[ "Cahill Center for Astronomy and Astrophysics 301-17, California Institute of Technology, Pasadena, CA 91125, USA; Sorbonne Universités, Université Pierre et Marie Curie, Paris 6, CNRS, Observatoire de Paris, UMR 8112, LERMA, Paris, France;", "I. Physikalisches Institut, University of Cologne, Zülpicher Strasse 77, D-50937 Köln, Germany", "University of Michigan, Ann Arbor, MI 48109, USA", "École Normale Superiéure, CNRS, Observatoire de Paris, UMR 8112, LERMA, Paris, France", "Department of Earth and Space Sciences, Chalmers University of Technology, Onsala Space Observatory, SE-43992 Onsala, Sweden", "I. Physikalisches Institut, University of Cologne, Zülpicher Strasse 77, D-50937 Köln, Germany", "École Normale Superiéure, CNRS, Observatoire de Paris, UMR 8112, LERMA, Paris, France", "École Normale Superiéure, CNRS, Observatoire de Paris, UMR 8112, LERMA, Paris, France", "I. Physikalisches Institut, University of Cologne, Zülpicher Strasse 77, D-50937 Köln, Germany", "École Normale Superiéure, CNRS, Observatoire de Paris, UMR 8112, LERMA, Paris, France", "Jet Propulsion Laboratory, California Institute of Technology, 4800 Oak Grove Drive, Pasadena, CA 91109, USA", "Department of Physics and Astronomy, University of Toledo, Toledo, OH 43606, USA", "Cahill Center for Astronomy and Astrophysics 301-17, California Institute of Technology, Pasadena, CA 91125, USA", "Jet Propulsion Laboratory, California Institute of Technology, 4800 Oak Grove Drive, Pasadena, CA 91109, USA" ]
[ "2014A&A...566A..29S", "2014FaDi..168....9V", "2015A&A...574A..56G", "2015ARA&A..53..199G", "2015ApJ...800...40I", "2015MNRAS.450.3191O", "2015MNRAS.454.2292O", "2016A&A...585A.105L", "2016A&A...590A.105B", "2016ARA&A..54..181G", "2016PhRvA..93e2506O", "2017IAUS..322...75K", "2018ApJ...869..121M", "2019ApJ...883...54O", "2020MNRAS.497.1521A", "2021A&A...651A...9M", "2021IJAsB..20..111J", "2021JMoSp.37911482O", "2022arXiv220410908J", "2023MNRAS.518.3593D" ]
[ "astronomy" ]
7
[ "astrochemistry", "galaxies: nuclei", "ISM: molecules", "molecular processes", "submillimeter: general", "techniques: spectroscopic", "Astrophysics - Astrophysics of Galaxies" ]
[ "1937ApJ....86..483S", "1940PASP...52..187M", "1951ApJ...113..441B", "1963Natur.200..829W", "1969Natur.221..626C", "1970ApJ...161L..43W", "1973ApJ...185..505H", "1979ApJ...234L.139L", "1983ARA&A..21..239H", "1985PhRvL..54.1787L", "1992CP....163...83O", "1993ApJ...407..606S", "1994A&A...289..579T", "1994ASPC...58..355B", "1995A&A...294..667H", "1996PASJ...48..249K", "1999ARA&A..37..311E", "2000ApJ...543..764J", "2002A&A...383..603C", "2002A&A...395..357C", "2005ARA&A..43..677S", "2005ApJ...632..882O", "2005SSRv..119...71H", "2006ARA&A..44..367S", "2007A&A...468..627V", "2007mcim.book.....F", "2008A&A...485..451H", "2008AJ....135.1301V", "2009ApJ...705.1548R", "2009SPIE.7215E..0LH", "2010A&A...518L...1P", "2010A&A...518L...6D", "2010A&A...518L.108N", "2010A&A...518L.110G", "2010A&A...521L..20B", "2010ApJ...719..143T", "2012A&A...537A..17R", "2012ApJ...754..105H", "2012ApJ...758...83I", "2012RSPTA.370.5162L", "2013A&A...550A..25G", "2013ApJ...772..105M", "2013ApJ...779...23M", "2013JPCA..117.9661L", "2013JPCA..117.9919G" ]
[ "10.1088/0004-637X/785/2/135", "10.48550/arXiv.1403.1207" ]
1403
1403.1207_arXiv.txt
The presence of molecules in interstellar space was predicted as early as 1930's (e.g., \citealt{swings37}), quickly followed by identification of simple diatomic species, CH, CN, CH$^+$ (e.g., \citealt{mckellar40}), in optical absorption spectra. These early observations lead to important insights into the physics and chemistry of the interstellar medium (e.g., \citealt{bates51}). However, the optical and UV observations are limited to diffuse or translucent clouds and applications of molecular spectroscopy to studies of the dense interstellar medium had to await the development of microwave detection techniques in the 1960s. Since the first microwave detections of molecules in space \citep{weinreb63,cheung69,wilson70}, molecular spectroscopy has been a premier tool for studying interstellar gas cloud physics, from the Milky Way \citep{evans99} to the high-redshift universe \citep{solomon05}. In order to use molecules as quantitative tracers of star-forming clouds, it is necessary to understand the radiative and collisional processes that govern the excitation of their internal quantum states \citep{flower07}. For decades, it has been theorized that, under interstellar conditions, some molecules form via highly-exothermic reactions between ions and neutral molecules \citep{herbst73}. However, quantitative models of chemistry and of excitation have conventionally been constructed independently, on the grounds that chemical time-scales are long compared with excitation time-scales. The most abundant interstellar molecule, molecular hydrogen, is a symmetric rotor without a permanent dipole moment, which only has weak quadrupole rovibrational transitions in its ground electronic state (e.g., \citealt{habart05}). Its UV electronic transitions can only be studied from space and the observations are limited to the relatively nearby diffuse and translucent clouds. The infrared rovibrational lines are accessible from the ground, but are only excited in warm, active environments, such as shocks or photon dominated regions. Consequently, trace molecules, with rich rotational spectra in the millimeter to far-infrared wavelength range have been used as proxies of H$_2$, to understand the physical conditions in the star-forming gas reservoir of the Milky Way and external galaxies. Of these, oblate symmetric top molecules, in particular ammonia \citep{ho83}, for which the $J=K$ level is the lowest energy level in each $K$ ladder, are of great utility for the temperature determination. The population of these ``metastable'' levels is assumed to be thermalized at the kinetic temperature of the medium, due to the absence of allowed radiative transitions out of these levels. The hydronium ion, \htop, is another oblate symmetric rotor, isoelectronic with ammonia, which like ammonia has the characteristic inversion splitting of its rotational levels. In the case of \htop, however, the splitting is very large \citep{liu85} and the inversion transitions occur at far-infrared wavelengths, as opposed to centimeter wavelengths for ammonia. The far-infrared \htop\ inversion transitions, up to (11,11), have only recently been detected for the first time by \emph{Herschel}, in absorption toward the central region of the Milky Way \citep{lis12}. Earlier HIFI observations of the hydronium ion toward Sagittarius~B2 indicated a high rotational temperature, \about 500~K, at velocities corresponding to the cloud envelope. \cite{lis12} suggested that, in regions exposed to ionizing radiation, such as those at the center of the Milky Way, highly-excited states of the hydronium ion and ammonia are populated mainly by exoergic chemical formation processes and temperatures derived from their spectra do not represent the physical temperature of the medium. The detection of highly-excited states of the hydronium ion in nearby active galaxies, with rotational temperatures similar to that observed in Sagittarius B2 \citep{gonzalez13}, demonstrated the universality of the shape of the observed population diagrams in \emph{active} environments. Here we present new sensitive \emph{Herschel} observations of the hydronium ion on the lines of sight toward Sagittarius B2(N) and W31C (G10.6-0.4), which clearly demonstrate that this effect is not limited to the environments of galactic nuclei, characterized by high X-ray and cosmic ray fluxes, but is a \emph{fundamental} property, pervasive throughout the interstellar medium.
Earlier ground-based \citep{hutte95} and space-borne \citep{ceccarelli02} observations of ammonia absorption toward Sagittarius B2 have been interpreted in terms of a physically hot component of molecular gas in the central region of the Galaxy, in which the observed high rotational excitation temperature traces a correspondingly high kinetic temperature. At first glance this seems plausible, because shock waves in this active high-mass star-forming region should be able to maintain some molecular gas at temperatures of 500 to 700~K, or because the region contains luminous heating sources of X-rays and cosmic rays. Moreover, strong X-ray emission in iron fluorescence at 6.4~keV has been taken as evidence that the envelope of Sagittarius B2 was recently illuminated by an X-ray flash, attributable to a flare from the region around the central black hole that is associated with the Sagittarius A$^*$ radio continuum source \citep{sunyaev93, koyama96}. The Sagittarius B2 X-ray flux is now fading \citep{terrier10}, with a characteristic decay time of 8 years, comparable to the light-travel time across the central part of the molecular cloud. While the competing mechanisms put forward to explain the hot gas component on the line of sight toward Sagittarius B2 are all plausible, a generally accepted explanation has yet to be offered. Interestingly, recent observations of the (8,8) to (15,15) ammonia inversion lines \citep{mills13} indicate rotational temperatures of 350--450 K in several largely \emph{quiescent} Galactic center giant molecular clouds. This indicates that the effect is not limited to Sagittarius B2, but widespread in the Galactic center environment. The hot gas component traced by the high-energy ammonia inversion lines would thus have to fill a few hundred parsec size region. The energy input required to heat such a vast volume of gas would then have to be carefully evaluated. Moreover, the diffuse, widespread, low-density gas component in the Central Molecular Zone can be independently probed through the infrared absorption spectroscopy of the H$_3^+$ rovibrational transition \citep{oka05}, which indicate temperatures of ``only'' \about 250~K. Inversion lines of ammonia, up to (10,10), have also been detected in absorption toward PKS~1830-211, indicating rotational temperatures \ga 600~K for the highest-energy lines \citep{henkel08}. The absorption is attributed to molecular gas in spiral arms of an ordinary galaxy at a redshift of 0.89. Explanations similar to those previously invoked to explain the high-rotational temperatures of ammonia cannot easily be applied to H$_3$O$^+$: if the density were high enough to thermalize the populations of highly excited rotational states at a kinetic temperature of 400 to 600~K, then molecular collisions would also raise the excitation temperatures of the high-$J$ inversion doublets to values higher than the brightness temperature of the continuum radiation at $\lambda 150 - 200$~\mic. In that case, the inversion lines could not appear in absorption as is observed. A fundamentally different interpretation is required to explain simultaneously the high rotational excitation of H$_3$O$^+$, its short chemical lifetime, and its relatively large total column density, $N({\rm H}_3{\rm O}^+) = 4-7\times 10^{14}$ cm$^{-2}$, close to an intense far-infrared continuum source. The characteristic two-component rotational diagram for metastable states of H$_3$O$^+$ naturally results from its formation by exoergic ion-neutral reactions. The excess enthalpy (heat of formation) of the main source reaction (e.g., \citealt{herbst73}) $$ {\rm H}_2{\rm O}^+ + {\rm H}_2 \to {\rm H}_3{\rm O}^+ + {\rm H} $$ goes partly into internal rotational excitation of the product ion, the highly excited non-metastable states relax rapidly by spontaneous radiative transitions to the metastable states (i.e. the lower $J=K$ inversion levels), whose populations are limited by the rates of excited-molecule formation and chemical destruction. Populations are partly re-distributed by inelastic collisions and by absorption of the background far-infrared continuum radiation. The populations of the lowest metastable states may approach a collision-dominated thermal distribution. The turnover to a ``hot'' formation-dominated distribution is detectable when the formation rate is a non-negligible fraction of the inelastic collision rates. The relation between molecular excitation and formation can be outlined quantitatively as follows. The destruction rate per H$_3$O$^+$ molecule by dissociative recombination with electrons can be written ${\cal D} = n(e) k_{\rm dr}$ s$^{-1}$, where the rate coefficient for the process $k_{\rm dr} = 4.3\times 10^{-7} (T/300)^{-1/2}$ cm$^3$ s$^{-1}$ at kinetic temperature $T$ \citep{jensen00} and $n(e)$ is the number density of free electrons. At $n(e)\sim 0.01$ to $0.1$ cm$^{-3}$, and $T\sim 50$ to $200$ K, ${\cal D} \approx 10^{-8}$ to $10^{-7}$ s$^{-1}$. This is fast enough to be the dominant destruction mechanism. If we furter assume steady state between destruction and formation, then the formation rate per unit voulme can be written as ${\cal F} = n({\rm H}_3{\rm O}^+) {\cal D}$~cm$^{-3}$s$^{-1}$. With no further knowledge of the chemistry, we can assert that the source of H$_3$O$^+$ in the envelope of Sagittarius B2 must be of the order of $$ {\cal F} = N({\rm H}_3{\rm O}^+) {\cal D} / L \approx 10^{-11} \Bigl[{{1 \;{\rm pc}}\over{L}}\Bigr] \Bigl[{{n(e)}\over{0.1\;{\rm cm}^{-3}}}\Bigr] \;{\rm cm}^{-3}{\rm s}^{-1}, $$ where $L$ is the characteristic length scale of the absorbing region. The interstellar chemistry of oxygen-containing ions is thought to be straightforward \citep{herbst73}, with cosmic-ray or X-ray ionizations of hydrogen leading to H$^+$ and H$_3^+$, which transfer charge to O$^+$ or OH$^+$. The oxygen ions form H$_2$O$^+$ and finally H$_3$O$^+$ via reactions with H$_2$. In regions where the atomic fraction is small, H/H$_2 \ll 1$, the chemical source rate of the terminal ion H$_3$O$^+$ can be a significant fraction of the hydrogen ionization rate. Thus we can express ${\cal F} \sim \zeta n_{\rm H} \epsilon$, where $\zeta$ is the ionization rate, $n_{\rm H}$ is the density of hydrogen, and $\epsilon$ is a chemical loss factor that accounts for minor channels in the ion chemistry that prevent the hydrogen ions from processing solely the oxygen. Such an analysis has been applied to Herschel/HIFI observations of OH$^+$ and H$_2$O$^+$ in diffuse molecular gas by \cite{neufeld10}, who concluded that those highly reactive ions arise mostly in components of gas with a high atomic content, unlike \htop, which traces predominantly molecular gas. Following these arguments, we see that the observed column density and estimated destruction rate of H$_3$O$^+$ toward Sagittarius B2 require $$ \zeta n_{\rm H} \leq 10^{-11}/\epsilon \;\;\;{\rm cm}^{-3}\;{\rm s}^{-1} \;,$$ for $L\approx 1$ pc and $n(e)\leq 0.1$ cm$^{-1}$ in the envelope of Sagittarius B2. This could be achieved with a cosmic-ray ionization rate $\zeta \sim 10^{-15}$ s$^{-1}$ at $n_{\rm H}\epsilon \sim 10^4$ cm$^{-3}$, or with less extreme values if the electron density is lower. Cosmic ray ionization rates in excess of $10^{-15}$ s$^{-1}$ have been deduced from infrared absorption observations of H$_3^+$ on several lines of sight through the Central Molecular Zone of the Galaxy (\citealt{gotto13}, and references therein). The line of sight to W31C presents a simpler test case because it lies entirely outside the Central Molecular Zone, where conditions are extreme. For illustration, consider the low-density limit, in which collisional excitation is neglected and the rotational populations are governed entirely by the formation process and radiative transitions. Let the state-specific formation rate of H$_3$O$^+$ be described by $$ F(J,K,p) = {\cal F} g(J,K,p) \exp\bigl(-E(J,K,p)/kT_f \bigr) / Q(T_f) $$ for each state of rotational quantum numbers $J$, $K$, and parity $p$, where $g$ is the statistical weight, $E(J,K,p)$ is the energy, $Q$ is the partition function, and $T_f$ is a parameter---the formation temperature. The populations of the vibration-rotation levels are then computed by solving the rate equations with formation, destruction, and all radiative transitions included (equations 12--15 of \citealt{vandertak07}), but with all the collisional terms set to zero. The results of one example are compared with the observed population diagram in Figure~4. In this case, ${\cal D}=10^{-7}$ s$^{-1}$ (corresponding to $n(e)\approx 0.1$ cm$^{-3}$ at $T<100$ K) and the total column density is $N({\rm H}_3{\rm O}^+) = 1.4\times 10^{13}$ cm$^{-2}$. The formation temperature is $T_f = 400$ K and the molecules are exposed to an average Galactic background continuum similar to that described by \cite{black94}. At the adopted destruction rate, ${\cal F} = 4.5\times 10^{-14}$ cm$^{-3}$ s$^{-1}$ is required to attain a good match in column density over a total path length $L\approx 10$ pc. This rate, in turn, would imply a rate of cosmic-ray ionization of H$_2$, $\zeta_2 = {\cal F}/n({\rm H}_2)\epsilon \approx 5\times 10^{-16}$ s$^{-1}$ for a density $n({\rm H}_2) = 300$ cm$^{-2}$ and $\epsilon\approx 0.3$. This ionization rate would be consistent with that derived by \cite{indriolo12} toward W51 IRS2 if the H$_3$O$^+$ is found preferentially in regions of higher molecular fraction and higher $\epsilon$ than the observed OH$^+$. These conditions are also in harmony with photochemical models of diffuse molecular gas as discussed by \cite{hollenbach12}. \begin{figure}[t] \centering \includegraphics[width=0.95\columnwidth,angle=0]{fig4.eps} \caption{Observed population diagram for \htop\ toward W31C (blue and red squares). Derived rotational temperatures of the warm and hot component are given together with the corresponding maximum uncertainties, computed assuming 35\% maximum uncertainties of the individual measurement. Green triangles show the \htop\ population of the metastable levels computed using the simple formation model described in the text. The model is included for illustrative purposes only and is not meant to be a fit to the data.} \label{fig:pop} \end{figure} A combined excitation-chemistry model, which includes explicitly the exothermicity of the chemical reactions leading to the formation of the hydronium ion and ammonia is beyond the scope of the present paper and will be presented separately (J. Black, in prep.) Such calculations are clearly needed to demonstrate that chemistry and excitation are closely coupled in some interstellar environments and should not be treated separately and to use \htop\ as a quantitative tracer of the cosmic ray ionization rate\footnote{One serious impediment is that the collisional cross-sections for \htop\ are only available for rotational levels with energies below \about 350~K, scaled from ammonia \citep{offer92}.}. This aspect is of particular interest for understanding the energetics of the nuclear regions of ultraluminous, star-forming galaxies. Highly-excited inversion lines of the hydronium ion have now been detected in the ultraluminous infrared galaxy Arp~220, as well as in NGC~4418, which harbors a deeply buried, compact AGN \citep{gonzalez13}. In these environments, the hydronium ion absorption traces a relatively low-density molecular gas with a large filling factor, which is directly influenced by the feedback process of the central engine. The traditional explanation for these observations invokes a new phase of the interstellar medium, which would have to be heated by some combinations of X-rays, cosmic rays, shocks, or turbulence to \about 500 K. If the chemical formation pumping explanation is correct, this additional energy input is not required---the observations can simply be explained by chemical formation pumping in a much lower temperature medium. This relaxes the heating requirements and simplifies the physical picture of the interstellar medium in these active environments. With the improved understanding of the excitation, the far-infrared inversion lines of the hydronium ion also provide an accurate measurement of the ionization rate in the mostly molecular gas component of high-redshift galaxies, using ALMA, highly complementary to the OH$^+$ and H$_2$O$^+$ observations that primarily trace atomic gas, with low molecular fraction \citep{gerin10, neufeld10}. Symmetric top molecules have been proven to be some of the best tracers of the gas kinetic temperature in the interstellar medium \citep{ho83, mangum13}. However, the results presented here suggest that they may sometimes fail as an interstellar thermometer, under specific conditions. This typically involves low-density, diffuse gas, as collisions and strong radiation field tend to drive the rotational level population to equilibrium conditions. Consequently, care has to be taken when interpreting the derived rotational temperatures as physical temperatures of the medium.
14
3
1403.1207
We present new Herschel observations of the (6,6) and (9,9) inversion transitions of the hydronium ion toward Sagittarius B2(N) and W31C. Sensitive observations toward Sagittarius B2(N) show that the high, ~500 K, rotational temperatures characterizing the population of the highly excited metastable H<SUB>3</SUB>O<SUP>+</SUP> rotational levels are present over a wide range of velocities corresponding to the Sagittarius B2 envelope, as well as the foreground gas clouds between the Sun and the source. Observations of the same lines toward W31C, a line of sight that does not intersect the Central Molecular Zone but instead traces quiescent gas in the Galactic disk, also imply a high rotational temperature of ~380 K, well in excess of the kinetic temperature of the diffuse Galactic interstellar medium. While it is plausible that some fraction of the molecular gas may be heated to such high temperatures in the active environment of the Galactic center, characterized by high X-ray and cosmic-ray fluxes, shocks, and high degree of turbulence, this is unlikely in the largely quiescent environment of the Galactic disk clouds. We suggest instead that the highly excited states of the hydronium ion are populated mainly by exoergic chemical formation processes and the temperature describing the rotational level population does not represent the physical temperature of the medium. The same arguments may be applicable to other symmetric top rotors, such as ammonia. This offers a simple explanation of the long-standing puzzle of the presence of a pervasive, hot molecular gas component in the central region of the Milky Way. Moreover, our observations suggest that this is a universal process not limited to the active environments associated with galactic nuclei.
false
[ "such high temperatures", "Galactic interstellar medium", "high degree", "quiescent gas", "Galactic", "a high rotational temperature", "Sagittarius B2(N", "galactic nuclei", "~380 K", "exoergic chemical formation processes", "Sagittarius", "high X-ray and cosmic-ray fluxes", "the Galactic disk clouds", "Sensitive observations", "the rotational level population", "Sun", "the high, ~500 K, rotational temperatures", "other symmetric top rotors", "the foreground gas clouds", "the kinetic temperature" ]
12.386053
9.486311
-1
769663
[ "Ruiz-Lapuente, Pilar" ]
2014NewAR..62...15R
[ "New approaches to SNe Ia progenitors" ]
65
[ "Instituto de Física Fundamental, Consejo Superior de Investigaciones Científicas, c/. Serrano 121, E-28006 Madrid, Spain" ]
[ "2014ApJ...789L..45M", "2014ApJ...792...71P", "2014MNRAS.444.2157L", "2014MNRAS.445.2535S", "2015A&A...576A..86W", "2015ApJ...809..183X", "2015IJMPD..2430029M", "2015MNRAS.447.2803L", "2015MNRAS.450.2948A", "2015NuPhA.937...17B", "2015arXiv151202608V", "2016A&A...588A..88M", "2016ApJ...820..142R", "2016ApJ...826...53A", "2016MNRAS.455.4087G", "2016MNRAS.457..525G", "2016MNRAS.460.1614V", "2016MNRAS.462.2486F", "2016PASP..128e1001S", "2016PASP..128h2001H", "2016PhRvD..94b5026W", "2016arXiv161007230Y", "2017ApJ...835...97B", "2017ApJ...842..112R", "2017MNRAS.464.1607Y", "2017MNRAS.464.3965W", "2017MNRAS.465.2060B", "2017MNRAS.472.1593W", "2017PhRvD..95d3006W", "2017hsn..book.1211S", "2017suex.book.....B", "2018ApJ...862..124R", "2018EPJA...54..221L", "2018MNRAS.473..336B", "2018RAA....18...49W", "2018SCPMA..61d9502S", "2019A&A...622A..35L", "2019ApJ...870..135R", "2019ApJ...871...31A", "2019ApJ...871..177G", "2019ApJ...886...99L", "2019MNRAS.483..628S", "2019NewAR..8501523R", "2019arXiv191200900K", "2019sros.confE..80G", "2020A&A...638A.131N", "2020ApJ...895...70S", "2020PhRvL.124n1802R", "2021IJMPD..3050034K", "2021NJPh...23c1201A", "2021PhRvL.126m1101H", "2021arXiv210703568H", "2021arXiv211209656D", "2021gacv.workE..30S", "2022ApJ...935L...2H", "2022MNRAS.511.5018A", "2022PhRvC.106d5803D", "2022RAA....22b5001C", "2023ApJ...947...90R", "2023FrASS..1012880R", "2023MNRAS.520.1362D", "2023RAA....23h2001L", "2023RNAAS...7..129D", "2023TJAA....4S...7N", "2024arXiv240216635P" ]
[ "astronomy" ]
7
[ "Astrophysics - High Energy Astrophysical Phenomena", "Astrophysics - Cosmology and Nongalactic Astrophysics", "Astrophysics - Solar and Stellar Astrophysics" ]
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[ "10.1016/j.newar.2014.08.002", "10.48550/arXiv.1403.4087" ]
1403
1403.4087_arXiv.txt
\label{} Type Ia supernovae (SNe Ia) have been the tool that made possible the discovery of the acceleration of the expansion of the Universe (Riess et al. 1998; Perlmutter et al. 1999), and they are now providing new insights on the cosmic component, dubbed ``dark energy'', thus revealed. However, in contrast with their key role as cosmological probes, and after more than 50 years of supernova research, the nature of their progenitors remains elusive. As far back as 1960, it was established that Type I supernovae (in fact, the now denominated SNe Ia, or thermonuclear supernovae) should result from the ignition of degenerate nuclear fuel in stellar material (Hoyle \& Fowler 1960). The absence of hydrogen in the spectra of the SNe Ia almost immediately suggested that they were due to thermonuclear explosions of white dwarfs (WDs). Isolated white dwarfs were once thought to be possible progenitors (Finzi \& Wolf 1967), but soon discarded due to incompatibility with basic results from stellar evolution. Instead, accretion of matter from a close companion star in a binary system, by a previously formed C+O white dwarf with a mass close to the Chandrasekhar mass, provides a viable mechanism to induce the explosion (Wheeler \& Hansen 1971). Two main competing production channels are still under discussion nowadays. One possible path is the so--called single degenerate (SD) channel, where a C+O white dwarf grows in mass by accretion from a non--degenerate stellar companion: a main sequence star, a subgiant, a helium star, a red giant, or an AGB star (Whelan \& Iben 1973; Nomoto 1982). Another possible path is the double degenerate (DD) channel (Webbink 1984; Iben \& Tutukov 1984), where two WDs merge due to the loss of angular momentum by gravitational radiation. The merging could produce the collapse of the white dwarf (Saio \& Nomoto 1985), or it can produce a larger C+O white dwarf configuration that then explodes (Pakmor et al. 2012). \bigskip In the decade of the 90's, the variety amongst SNe Ia was discovered, ranging from events such as SN 1991bg to those as SN 1991T, through normal SNe Ia (see Filippenko 1997a,b; Branch et al. 2007; Leibundgut 2011). Such diversity was made amenable for cosmology when the correlation of the luminosity at the maximum of the light curve of each SN Ia with its rate of decline was parameterized (Phillips 1993, 1999; Riess, Press \& Kirshner 1995; Perlmutter et al. 1997). It became clear, then, that SNe Ia could be used as distance indicators in cosmology, and that led to the aforementioned discovery. Yet, the first decade of the new century has brought new surprises: super--Chandrasekhar supernovae, as well as extremly faint ones (see below). Neither of them are useful for cosmology, although they are not a severe nuisance there, since they can be easily identified, and eliminated from the large samples of SNe Ia collected for cosmological probes. Also, various teams have started to measure supernova rates at a wide variety of redshifts. The idea of using SNe Ia rates to discover the nature of the progenitor systems has now become an active line of research. Finally, high--resolution spectroscopic observations of SN have yielded the surprising result of time--varying absorptions, which indicate the existence of outflows in the circumstellar medium surrounding some SN, and points to possible nova activity previous to the explosion. An intriguing C II feature has been identifieed, close to the Si II line typical of SNe Ia, and that has led to thinking in two different directions: either the thermonuclear flame does not burn the outermost layers of the white dwarf, or maybe C is a signature of the merged white dwarf companion of the SN. There are also better estimates of the maximum H mass that could be present in the envelopes of the pre--SNe, if the explosions were triggered by accretion from a non--degenerate companion. There is continued failure to detect H from the radio emission of the SNe Ia, and there could be constraints from the X--ray emission as well. The task of searching for the companion star in Galactic supernovae has already given some definite results, and there are, now, simulations of the impact of the SN ejecta on the companion star that can be compared with the observations. In the following Sections, we present and discuss those new results. In Section 2 we briefly review the different models proposed to explain the SN Ia phenomenon. Section 3 examines how the Delay Time Distribution (DTD) constrains the possible SN Ia progenitors. In Section 4 we discuss the carbon and oxygen absorption features seen, in recent years, in the spectra of SN Ia at early times, while Section 5 deals with the emission features at late times. Section 6 discusses the variable blueshifted sodium feature seen in some SNe Ia. The X--ray constraints are presented in Section 7, and the radio constraints in Section 8. In Section 9 we report the limits on the luminosities of the companions of SNe Ia obtained from pre--explosion images. Section 10 deals with the detection of companions throught the early light curves of SNe Ia. Section 11 reviews the direct searches for surviving companions, in the Galaxy and in the Large Magellanic Cloud. Section 12 deals with the identification of possible candidates to SNe Ia through reconstruction of the orbital evolution of diverse close binary systems containing white dwarfs. Section 13 addresses the important problem of the outliers from the peak brightness--decline rate of the light curve relationship used to make these SNe calibrated candles for cosmology. Section 14 deals with the bulk of SNe Ia used for cosmology. We summarize the current state of affairs in the last Section.
14
3
1403.4087
Although Type Ia supernovae (SNe Ia) are a major tool in cosmology and play a key role in the chemical evolution of galaxies, the nature of their progenitor systems (apart from the fact that they must content at least one white dwarf, that explodes) remains largely unknown. In the last decade, considerable efforts have been made, both observationally and theoretically, to solve this problem. Observations have, however, revealed a previously unsuspected variety of events, ranging from very underluminous outbursts to clearly overluminous ones, and spanning a range well outside the peak luminosity-decline rate of the light curve relationship, used to make calibrated candles of the SNe Ia. On the theoretical side, new explosion scenarios, such as violent mergings of pairs of white dwarfs, have been explored. We review those recent developments, emphasizing the new observational findings, but also trying to tie them to the different scenarios and explosion mechanisms proposed thus far.
false
[ "white dwarfs", "new explosion scenarios", "explosion mechanisms", "calibrated candles", "SNe Ia", "violent mergings", "pairs", "galaxies", "the light curve relationship", "Type Ia", "considerable efforts", "the new observational findings", "cosmology", "the different scenarios", "the SNe Ia.", "the peak luminosity-decline rate", "their progenitor systems", "events", "the theoretical side", "the chemical evolution" ]
2.948622
5.321352
21
12308915
[ "Chaty, Sylvain" ]
2013AdSpR..52.2132C
[ "Optical/infrared observations unveiling the formation, nature and evolution of High-Mass X-ray Binaries" ]
33
[ "-" ]
[ "2014A&A...568A..76D", "2014ApJ...797..114S", "2014arXiv1412.4699B", "2015A&A...577A.130F", "2015A&A...578A.107M", "2015AdSpR..55.1255B", "2015ApJ...808...34P", "2015arXiv150103882T", "2015arXiv150301087G", "2016A&A...591A..87C", "2017ApJ...835..249Z", "2017ApJ...841...35F", "2017ApJ...846..170T", "2018A&A...618A.150F", "2018ApJ...856...38K", "2018mgm..conf.1883C", "2019A&A...628A..19Q", "2019A&A...631A.135D", "2019IAUS..346...49F", "2019IAUS..346..161C", "2019IAUS..346..212G", "2020A&A...637A...2C", "2020ApJ...894...86F", "2021A&A...647A..75F", "2021MNRAS.501.1116A", "2022A&A...659A..48S", "2022A&A...665A..31F", "2022MNRAS.517.3034B", "2022abn..book.....C", "2023A&A...671A.149F", "2023MNRAS.521.4039S", "2023arXiv231115778C", "2023pbse.book.....T" ]
[ "astronomy" ]
7
[ "Astrophysics - High Energy Astrophysical Phenomena", "Astrophysics - Solar and Stellar Astrophysics" ]
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[ "10.1016/j.asr.2013.09.010", "10.48550/arXiv.1403.0792" ]
1403
1403.0792.txt
Nearly fifty years after the discovery of the first extra-solar X-ray source \citep[Sco\,X-1;][]{giacconi:1962}, X-ray astronomy has now reached its age of reason, with a plethora of telescopes and satellites covering the whole electromagnetic spectrum, obtaining precious observations of these powerful celestial objects populating our Universe. %\subsection{Different types of X-ray binaries} High energy binary systems are composed of a compact object -- neutron star (NS) or black hole (BH) -- orbiting a companion star, and accreting matter from it \citep[cf e.g.][for a review]{chaty:2008a}. The companion star is either a low-mass star (typically $\sim 1 \Msol$ or less, with a spectral type later than B, called in the following LMXB for ``Low-Mass X-ray Binary''), or a luminous early spectral type OB high-mass companion star (typically $> 10 \Msol$, called HMXB for ``High-Mass X-ray Binary''). 300 high energy binary systems are known in our Galaxy: 187 LMXBs and 114 HMXBs \citep[respectively 62\% and 38\% of the total number;][]{liu:2006,liu:2007}. Accretion of matter is different for both types of sources\footnote{We will not review here the so-called $\gamma$-ray binaries, where most of the energy is in the GeV-TeV range, coming from interaction between relativistic electrons of the pulsar and the wind of the massive star.}. In the case of LMXBs, the small and low-mass companion star fills and overflows its Roche lobe, therefore accretion of matter always occurs through the formation of an accretion disc. The compact object can be either a NS or a BH, Sco X-1 falling in the former category. % In the case of HMXBs, accretion can also occur through an accretion disc, for systems in which the companion star overflows its Roche lobe; however this is generally not the case, and there are two alternatives. The first one concerns stars with a circumstellar disc, and here it is when the compact object -- on a wide and eccentric orbit -- crosses this disc, that accretion periodically occurs (case of HMXBs containing a main sequence early spectral type Be III/IV/V star, rapidly rotating, called in the following BeHMXBs). The second case is when the massive star ejects a slow and dense wind radially outflowing from the equator, and the compact object directly accretes the stellar wind through e.g. Bondy-Hoyle-Littleton processes (case of HMXBs containing a supergiant I/II star, later called sgHMXBs). %(cf Figure \ref{lmxb-hmxb}: LMXB on the left and HMXB on the right). We point out that there also exists binary systems for which the companion star possesses an intermediate mass (typically between 1 and $10 \Msol$, called IMXB for "Intermediate-Mass X-ray Binary"). We first describe in Section \ref{section:HMXBs} the various existing types of HMXBs, we then report in Section \ref{section:INTEGRAL} the new populations of HMXBs discovered by the {\it INTEGRAL} satellite, we give further details in Section \ref{section:formation-evolution} on the formation and evolution of HMXBs, and we finally conclude this review in Section \ref{section:conclusion}.
\label{section:conclusion} While the {\it INTEGRAL} satellite was not primarily designed for this, it allowed a great progress in the understanding of HMXBs in general, and of sgHMXBs in particular. Let us first recall the {\it INTEGRAL} legacy concerning sgHMXBs: \begin{itemize} \item The {\it INTEGRAL} satellite has quintupled the total number of known Galactic sgHMXBs, constituted of a NS orbiting a supergiant star. Most of the new sources are slow and absorbed X-ray pulsars, exhibiting a large $\nh$ and long $P_\mathrm{spin}$ ($\sim1$\,ks), typical of wind-fed accreting pulsars (according for instance to the Corbet diagram). \item The {\it INTEGRAL} satellite has revealed the existence in our Galaxy of two previously hidden populations of high energy binary systems. First, a population of obscured and persistent sgHMXBs, composed of supergiant companion stars exhibiting a strong intrinsic absorption and long $P_\mathrm{spin}$, with the NS deeply embedded in the dense stellar wind, forming a dust cocoon enshrouding the whole binary system. Second, the SFXTs, exhibiting brief and intense X-ray flares -- with a {\bf luminosity $L_X \sim 10^{36} \ergs$ at the peak, during a few ks} every $\sim 7$\,days --, which can be explained by accretion through clumpy winds. \end{itemize} However, one may ask the questions: Apart from these observational facts, has the {\it INTEGRAL} satellite allowed us to better understand all populations of HMXBs, by increasing their number, compared to the ones already known before its launch? Do we better apprehend the accretion processes in HMXBs in general, and in sgHMXBs in particular, and what makes the fast transient flares so special, in the context of the clumpy wind model, of the formation of transient accretion discs, and/or the centrifugal/magnetic barrier? The answers to these questions is probably {\it ``not yet''}, however we now have in hand more sources, and therefore more constraints to play with. Studying these populations will provide a better understanding of the formation and evolution of short-living HMXBs, and study accretion processes. Finally, it is clear that stellar population synthesis models now have to take these objects into account, to assess a realistic number of all populations of high energy binary systems in our Galaxy. %
14
3
1403.0792
In this review I first describe the nature of the three kinds of High-Mass X-ray Binaries (HMXBs), accreting through: (i) Be circumstellar disc, (ii) supergiant stellar wind, and (iii) Roche lobe filling supergiants.
false
[ "filling supergiants", "circumstellar disc", "Roche", "High-Mass X-ray Binaries", "i", "HMXBs", "Binaries", "(ii) supergiant stellar wind", "(iii) Roche", "the nature", "the three kinds", "this review", "I", "first", "three" ]
6.540584
6.367085
-1
813198
[ "Brax, Philippe", "Valageas, Patrick" ]
2014PhRvD..90b3508B
[ "K-mouflage cosmology: Formation of large-scale structures" ]
47
[ "Institut de Physique Théorique, CEA, IPhT, 91191 Gif-sur-Yvette, France and CNRS, URA 2306, 91191 Gif-sur-Yvette, France", "Institut de Physique Théorique, CEA, IPhT, 91191 Gif-sur-Yvette, France and CNRS, URA 2306, 91191 Gif-sur-Yvette, France" ]
[ "2014JCAP...09..031B", "2014PhRvD..90l3521B", "2015JCAP...08..028B", "2015MNRAS.454.2722H", "2015MNRAS.454.4085B", "2015PhRvD..91f3528B", "2015PhRvD..91j3503B", "2015PhRvD..91l3507W", "2015PhRvD..91l3522B", "2015PhRvD..92d3519B", "2016JCAP...01..020B", "2016JCAP...04..032T", "2016JCAP...05..001B", "2016JCAP...06..033G", "2016JCAP...11..039L", "2016MNRAS.456..146L", "2016PhRvD..93j3524P", "2016PhRvD..94d3529B", "2016PhRvD..94d4027B", "2017JCAP...02..031B", "2017JCAP...02..050B", "2017MNRAS.469..705S", "2017PhRvD..95d3515B", "2017PhRvD..96l4009R", "2018IJMPD..2748009B", "2018JPhCS.956a2001B", "2018MNRAS.481.2813G", "2018RPPh...81a6902B", "2019JCAP...01..049B", "2019JCAP...05..027B", "2019MNRAS.484.1149P", "2019PhRvD.100d3537D", "2019arXiv190803430B", "2020JCAP...05..002C", "2020MNRAS.497.1885H", "2020PhRvD.101f4065E", "2020arXiv200413677B", "2021JCAP...11..012D", "2021PhRvD.103j3505J", "2021RvMP...93a5003B", "2022JCAP...01..048H", "2022JCAP...05..030S", "2023MNRAS.519.4780B", "2023Univ....9..302H", "2023arXiv230612368F", "2024arXiv240612576B", "2024arXiv240613667B" ]
[ "astronomy", "physics" ]
10
[ "98.80.-k", "Cosmology", "Astrophysics - Cosmology and Nongalactic Astrophysics", "General Relativity and Quantum Cosmology" ]
[ "1967ApJ...147...73S", "1997ApJ...490..493N", "2002PhRvD..65j3510C", "2004AIPC..736..105B", "2004PhRvD..70h3536A", "2004Prama..63..731W", "2006IJMPD..15.1753C", "2009A&A...508...93V", "2009ApJ...701..414G", "2009IJMPD..18.2147B", "2011PhRvD..83j4026B", "2012ApJ...744....3M", "2012JCAP...10..002B", "2012PhLB..715...38B", "2012PhRvD..86f3512B", "2013ApJS..208...19H", "2013JCAP...01..020B", "2013JCAP...04..029B", "2013JCAP...05..023L", "2013JCAP...10..027B", "2013JCAP...11..012L", "2013LRR....16....6A", "2013MNRAS.428..743L", "2013PhRvD..87h3522V", "2013PhRvD..88b3527B", "2013PhRvD..88h3524V", "2013arXiv1312.2006K", "2014A&A...571A..16P", "2014PhRvD..90b3507B" ]
[ "10.1103/PhysRevD.90.023508", "10.48550/arXiv.1403.5424" ]
1403
1403.5424_arXiv.txt
\label{Introduction} Scalar fields could be crucial in explaining the recent acceleration of the expansion of the Universe \cite{Copeland:2006wr}. They could also modify gravity as described by General Relativity (GR) \cite{Khoury:2010xi}. Such scalar fields with low masses could affect the growth of structures on very large scales in the Universe. On the other hand, in the Solar System or the laboratory, modifications of General Relativity are tightly constrained \cite{Will:2004nx}. The compatibility between the two behaviors on large and small scales can be ascertained in screened modified gravity where environmental effects take place in the presence of matter. In this paper, we study the formation of structure in K-mouflage models \cite{Babichev:2009ee, Brax:2012jr}. The background cosmology has been analyzed in a companion paper \cite{Brax:2014aa}. The K-mouflage mechanism is present in scalar field theories with noncanonical kinetic terms. It is effective in regions of space where the gravitational acceleration is larger than a critical value \cite{Khoury:2013tda}. On large scales where matter is less dense, deviations from GR can be significant and affect the growth of structure. In particular, K-mouflage models do not converge towards GR in the large distance regime. As a result, K-mouflage models behave like a linear theory with a time dependent Newton constant up to quasilinear scales. We study the perturbative regime of K-mouflage models before analyzing the nonlinear regime. Linear perturbation theory differs from $\Lambda$-CDM on large scales and is therefore amenable to clean comparisons with the measurements of the growth factor as forecast by the EUCLID mission \cite{Amendola:2012ys}. It turns out that the linear regime in the scalar sector gives a good description of the growth structure up to quasilinear scales owing to the relative irrelevance of nonlinear corrections in the scalar sector. In the nonlinear regime, we use the spherical collapse to deduce the halo mass function and the deviation of the power spectrum from $\Lambda$-CDM. As expected the halo mass function is significantly affected for large masses while the power spectrum can see large deviations on nonlinear scales of order $1\ {\rm Mpc}$. This is a feature of $f(R)$ models too \cite{Brax2013} which could be disentangled here inasmuch as it persists even for moderate redshifts in the K-mouflage case. In fact, the redshift dependence of the deviations from GR is also very different in models with the Vainshtein property such as Galileons \cite{Barreira:2013eea, Li:2013tda}, with fewer deviations at moderate redshifts than in the K-mouflage models. Hence we can expect that the three screening mechanisms could be disentangled by both analyzing the large-scale features, as chameleonlike models converge to GR contrary to K-mouflage and Vainshtein models, and the time evolution of their deviations from GR, as K-mouflage models show persistent ones up to $z=2$. In section II, we introduce K-mouflage models, screening and the background cosmology. In section III, we analyze the perturbative regime of K-mouflage theories. In section IV, we focus on large scales and the ISW effect. In section V, we study the spherical collapse and apply these results to a halo model which allows us to tackle the cosmology of nonlinear scales. Finally in section VI, we calculate the power spectrum including nonlinearities as defined by a halo model. In section VII, we compare our results with chameleonlike models and Galileons. We conclude in section VIII. Two appendices present the perturbation theory of K-mouflage models and a comparison of the power spectrum in K-mouflage models with the one where the only modification from $\Lambda$-CDM is due to the K-mouflage background cosmology.
\subsection{Summary} Before we conclude this paper, let us briefly summarize the main properties of the K-mouflage models studied here and our results: - K-mouflage models involve an additional scalar field, $\varphi$, with a nonstandard nonlinear kinetic term, ${\cal M}^4 K[-(\pl\varphi)^2/2 {\cal M}^4]$, where ${\cal M}^4 $ is of the order of the critical density now. Here we consider models where the field $\varphi$ is also conformally coupled to matter fields through the Jordan metric $\tilde{g}_{\mu\nu} = A^2(\varphi) g_{\mu\nu}$. - The nonlinearity of the Lagrangian, which gives rise to terms $\bar{K}' (\pl\delta\varphi)^2$ for the fluctuations with respect to the cosmological background, provides a ``screening mechanism'' as the large prefactor $\bar{K}'$ freezes the fluctuations and suppresses the fifth force in the high-density and small-scale regimes. This provides a convergence to GR on small astrophysical scales (studied in an upcoming paper) and at high redshift \cite{Brax:2014aa}. On the other hand, in contrast with some other modified-gravity scenarios (e.g., chameleon and Galileon models), linear cosmological structures are unscreened and show deviations from $\Lambda$-CDM up to the Hubble scale. Moreover, the dark energy background evolution of K-mouflage models only behaves as a cosmological constant at low redshifts. - The equations of motion obtained for matter on cosmological scales are the usual continuity equation, a modified Euler equation (with an additional friction term and an additional fifth-force potential term), and a modified Poisson equation (with a time dependent effective Newton constant). The scalar field fluctuations obey a time dependent Klein-Gordon equation, as the background field evolves with time. - Even though this background does not follow a quasistatic evolution, on small scales (far below the horizon) the scalar field fluctuations obey a quasistatic regime (because spatial gradients dominate over time derivatives) and are ``slaved'' to the same-time density fluctuations (i.e., the Klein-Gordon equation takes the form of a nonlinear Poisson equation). - For cosmological structures, from the cosmic web down to clusters of galaxies, which show moderate density contrasts ($\delta \lesssim 200$), this quasistatic Klein-Gordon equation can be linearized. Therefore, these models provide an explicit nonlinear example where the scalar field sector can be linearized on cosmological scales (while the nonlinearity appears on smaller astrophysical scale and ensures the convergence back to GR). - Quasilinear scales can be studied using cosmological perturbation theory as in the standard $\Lambda$-CDM scenario, taking into account the new linear friction and fifth-force terms in the Euler equation. In particular, the nonlinearities are due to the usual transport terms that are identical to those found in the $\Lambda$-CDM case. - The linear regime growth factors differ from the $\Lambda$-CDM predictions through time dependent terms, but in contrast with some other modified-gravity models, they do not show an additional scale dependence (because large scales remain unscreened up to the horizon). The sign of the deviation from $\Lambda$-CDM is set by the sign of the derivative $\bar{K}'$, as for background quantities \cite{Brax:2014aa}. For instance, models with $K'>0$ yield a smaller Hubble expansion rate $H(z)$ (with a common normalization today) and larger linear growth rates $D_+(z)$ and $f(z)$. More precisely, the quantity that governs the deviations from the $\Lambda$-CDM predictions, both for the background and the perturbations, is the ratio $\beta^2/\bar{K}'$, where $\beta$ is the coupling constant to the matter. As for small-scale screening, these deviations are suppressed in models with a large nonlinear factor $\bar{K}'$. - Because large linear scales deviate from the $\Lambda$-CDM predictions, large-scale CMB anisotropies also show a significant deviation through the ISW effect. In particular, the cross-correlation between the large-scale CMB temperature fluctuations and low-redshift galaxy surveys can change sign for models with $K'<0$, and the amplitude of the relative deviation from $\Lambda$-CDM grows with redshift. This also gives rise to a deviation for the low-$\ell$ CMB multipoles $C_{\ell}$, but this generically yields more power than the $\Lambda$-CDM prediction over $\ell \leq 10$ whatever the sign of $\bar{K}'$. - To go beyond the perturbative regime, we have also studied the spherical collapse dynamics. For the same reason as the absence of scale dependence in the linear regime, the spherical collapse is only modified by time dependent but scale-independent factors. This also means that, as in GR or Newtonian gravity, different mass shells remain uncoupled until shell crossing. This simplifies the analysis and it leads to a time dependent linear density contrast threshold $\delta_{L}(z)$ for the collapse of virialized halos. This yields in turn a deviation for the large-mass tail of the halo mass function, that again depends on the sign of $\bar{K}'$. - Combining perturbation theory and the spherical collapse dynamics, we have estimated the matter power spectrum up to mildly nonlinear scales $k \lesssim 10 h$Mpc$^{-1}$. We recover a constant relative deviation from $\Lambda$-CDM on linear scales and a peak on weakly nonlinear scales, $k \sim 1 h$Mpc$^{-1}$, due to the amplification associated with the nonlinear matter dynamics and the large-mass tail of the halo mass function. Therefore, large-scale structures provide a useful probe of such models as the deviations from $\Lambda$-CDM for $P(k)$, or the halo mass function, can be greater by a factor of 10 than those of background quantities, such as $H(z)$. - These deviations decrease rather slowly at higher redshift and remain non-negligible at $z=2$ (as compared to $z=0$). This is due to the fact that the dark energy component only slowly becomes subdominant at high $z$, because its energy density actually grows (but at a smaller rate than the matter density). This feature is rather different from the behavior obtained in some other modified-gravity models [e.g., $f(R)$ theories or dilaton models] where the background is almost identical to $\Lambda$-CDM and the deviations for matter perturbations are only significant at low $z$. \subsection{Conclusions} In conclusion, K-mouflage is an alternative to the screening by the chameleon or the Vainshtein mechanisms with striking features on the growth of large-scale structures. The most significant one is certainly the absence of screening of large astrophysical objects on cosmological scales, such as galaxy clusters. In this regime, the scalar theory behaves like a linear field theory leading to a time dependent modification of Newton's constant and an increase/decrease of the growth of structure compared to $\Lambda$-CDM depending on the ratio $\beta^2/\bar K'$ corresponding to the square of the effective coupling to matter when the bare coupling $\beta$ is rescaled by the wave function normalization of the field $|\bar K'|^{1/2}$. For models where deviations from the $\Lambda$-CDM behavior at the background level are at the percent level, the deviations of the power spectrum of the density contrast on mildly nonlinear scales is enhanced compared to the linear part of the spectrum and can reach ten percent. Moreover, the convergence to the Einstein- de Sitter behavior of perturbations in the past is rather slow due to the properties of the background cosmology. Indeed, at the background level, the Hubble rate converges to the Einstein- de Sitter case in the distant past due to the screening of the scalar field in the high-density environment of the early Universe while it converges to a $\Lambda$-CDM behavior in the very recent past. In the intermediate regime around redshifts of $1\lesssim z \lesssim 2$, the Hubble rate can differ significantly from its $\Lambda$-CDM counterpart. This translates into a relative persistence of the deviations from $\Lambda$-CDM which differs from other screening mechanisms, up to redshifts of a few. We leave a more detailed analysis of cosmological observational constraints to future works. K-mouflage could also have different features on smaller scales where the density contrast is larger than in galaxy clusters. In this regime, the nonlinearities of the models reappear and cannot be neglected, especially on scales of the order of the K-mouflage radius. This is left for future work.
14
3
1403.5424
We study structure formation in K-mouflage cosmology whose main feature is the absence of screening effect on quasilinear scales. We show that the growth of structure at the linear level is affected by both a new time dependent Newton constant and a friction term which depend on the background evolution. These combine with the modified background evolution to change the growth rate by up to ten percent since z∼2. At the one loop level, we find that the nonlinearities of the K-mouflage models are mostly due to the matter dynamics and that the scalar perturbations can be treated at tree level. We also study the spherical collapse in K-mouflage models and show that the critical density contrast deviates from its Λ-CDM value and that, as a result, the halo mass function is modified for large masses by an order one factor. Finally we consider the deviation of the matter spectrum from Λ-CDM on nonlinear scales where a halo model is utilized. We find that the discrepancy peaks around 1 h Mpc<SUP>-1</SUP> with a relative difference which can reach fifty percent. Importantly, these features are still true at larger redshifts, contrary to models of the chameleon-f(R) and Galileon types.
false
[ "z∼2", "tree level", "models", "quasilinear scales", "Galileon types", "nonlinear scales", "large masses", "larger redshifts", "K-mouflage models", "the modified background evolution", "a halo model", "the halo mass function", "Newton", "Galileon", "the background evolution", "structure formation", "K-mouflage cosmology", "the growth rate", "the K-mouflage models", "the linear level" ]
10.060942
1.350501
-1
482675
[ "Mann, Rita K.", "Di Francesco, James", "Johnstone, Doug", "Andrews, Sean M.", "Williams, Jonathan P.", "Bally, John", "Ricci, Luca", "Hughes, A. Meredith", "Matthews, Brenda C." ]
2014ApJ...784...82M
[ "ALMA Observations of the Orion Proplyds" ]
109
[ "National Research Council Canada, 5071 West Saanich Road, Victoria, BC, V9E 2E7, Canada", "National Research Council Canada, 5071 West Saanich Road, Victoria, BC, V9E 2E7, Canada; Department of Physics and Astronomy, University of Victoria, Victoria, BC, V8P 1A1, Canada", "National Research Council Canada, 5071 West Saanich Road, Victoria, BC, V9E 2E7, Canada; Department of Physics and Astronomy, University of Victoria, Victoria, BC, V8P 1A1, Canada; Joint Astronomy Centre, 660 North A'ohoku Place, University Park, Hilo, HI 96720, USA", "Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138, USA", "Institute for Astronomy, University of Hawaii, 2680 Woodlawn Drive, Honolulu, HI 96822 USA", "CASA, University of Colorado, CB 389, Boulder, CO 80309, USA", "Department of Astronomy, California Institute of Technology, MC 249-17, Pasadena, CA 91125, USA", "Van Vleck Observatory, Astronomy Department, Wesleyan University, 96 Foss Hill Drive, Middletown, CT 06459, USA", "National Research Council Canada, 5071 West Saanich Road, Victoria, BC, V9E 2E7, Canada; Department of Physics and Astronomy, University of Victoria, Victoria, BC, V8P 1A1, Canada" ]
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[ "astronomy" ]
16
[ "protoplanetary disks", "submillimeter: planetary systems", "Astrophysics - Solar and Stellar Astrophysics" ]
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[ "10.1088/0004-637X/784/1/82", "10.48550/arXiv.1403.2026" ]
1403
1403.2026_arXiv.txt
\label{sec: intro} Circumstellar disks are the birthsites of exoplanets. {\it Hubble Space Telescope} ({\it HST}) images of the Orion Nebula Cluster (ONC) revealed a hostile environment; many of the disks that orbit low mass stars are being photoevaporated by the intense UV radiation from the most massive nearby star, $\theta^1$\,Ori C \citep[spectral type O6;][]{odell94, mccullough,bally98,smith,ricci}. These disks are surrounded by tear-drop shaped structures, with bright heads facing $\theta^1$\,Ori C and tails pointing radially away. These distinctive circumstellar morphologies led to the nomenclature ``proplyds", an acronym for PROtoPLanetarY DiskS, that is now regularly applied to low-mass stars and their disks in the centers of massive star forming regions \citep{odell94}. The Orion proplyds were found to suffer photoevaporative mass-loss rates of $\dot{M} \approx 10^{-7}$\,$M_\odot$\,yr$^{-1}$ \citep{churchwell,henney}, high enough to disperse the amount of disk mass required to form a planetary system like our own in under 1\,Myr. That dissipation timescale is too short compared to the core accretion requirements for giant planet formation \citep[e.g.,][]{hubickyj05}, and is in apparent conflict with the inferred ages of the ONC stars \citep[$\sim$2\,Myr;][]{reggiani11,dario}. Measurements of the masses that remain in the Orion proplyds are crucial for characterizing the photoevaporation process and assessing their potential for planet formation. The most straightforward way to estimate a disk mass is from a measurement of the thermal dust continuum luminosity at long wavelengths, where the emission is optically thin \citep[cf.][]{beckwith90}. Although molecular gas likely comprises the vast majority of the mass budget in disks, the dust dominates the opacity and is significantly easier to detect. The Berkeley-Illinois-Maryland Array (BIMA), Owens Valley Radio Observatory (OVRO), Plateau de Bure Interferometer (PdBI), and Combined Array for Research in Millimeter Astronomy (CARMA) observed the Orion proplyds at wavelengths of 1.3--3.5\,mm \citep{mundy,bally98b,elada98,eisner06,eisner08}, but unfortunately these data provided limited constraints on the disk masses: contamination by free-free radiation from the ionized cocoons generated by the photoevaporation process often dominated the dust emission. Soon after the Submillimeter Array (SMA) was commissioned, it produced the first successful detections of the Orion proplyds at a submillimeter wavelength (880\,$\mu$m), where the dust emission dominates \citep{williams}. Those observations revealed that at least some of these disks still have sufficient mass ($>$ 10\,M$_{\rm jup}$) remaining to potentially form giant planets. A larger scale SMA survey of the Orion proplyds identified the erosion of the high end of the disk mass distribution due to photoevaporation by $\theta^1$\,Ori C \citep{mann09a, mann10}. To date, however, these surveys have only been sensitive enough to detect the most massive disks ($\ge$\,8.4\,M$_{\rm jup}$) in the ONC, and therefore provide relatively biased information about disk evolution in the hearts of massive star forming regions. To probe the full disk mass distribution in the ONC and further study the impact of external photoionizing radiation on disk properties, we carried out a much more sensitive pilot survey with the Atacama Large Millimeter/submillimeter Array (ALMA) that targeted 48 young stars, including 22 {\it HST}-identified proplyds. In this article, we present the results of the 856\,$\mu$m (ALMA Band 7) continuum observations. The observations and data reduction are described in \S \ref{obs}. Estimates of disk masses are presented in \S \ref{results}, and an examination of the dependence of disk mass on location in the ONC is discussed in \S \ref{disc}.
\label{disc} We observed the 856\,$\mu$m continuum emission toward 48 young stars in the Orion Nebula cluster using ALMA in Cycle 0, including 22 {\it HST}-identified proplyds. With an overall 3$\sigma$ survey sensitivity limit of $\sim$1.2\,$M_{\rm jup}$, we detected 23 disks ($\sim$48\%), including 9 that had not been detected previously. Aside from the disks around 253--1536b and 113-438 (see Table \ref{table3}), these detections coincide with the optically discovered disks from {\it HST} observations, highlighting the sensitivity of the space telescope to ONC disks due to their contrast with the bright nebular background. The giant silhouette disk 114--426 was detected for the first time, and has a low flux density of 7\,mJy, and an estimated mass of 3.9\,$M_{\rm jup}$ (this disk will be the subject of a forthcoming article; J. Bally et al., in preparation). Using this ALMA survey and the results of previous observations with the SMA, we find clear, statistically significant evidence for a marked decrease in the 856$\micron$ disk luminosities of the Orion proplyds that have smaller projected separations from the massive star $\theta^1$\,Ori C. In the assumption that the emission is optically thin, and the dust temperature and opacity are the same for all the disks, this implies that the masses of the Orion proplyds decrease for those disks located near $\theta^1$\,Ori C. The origins of that latter relationship could potentially be due to projection artifacts, initial conditions, and/or real evolutionary effects. The true separations between the proplyds and $\theta^1$\,Ori C are not known, but one can make a probabilistic argument that relates the projected separations to the true ones for an assumed distribution of orbital eccentricities around the ONC center of mass \citep[cf.][]{torres99}. For a uniform eccentricity distribution, the projected and true separations should be commensurate within a factor of $\sim$2; for a steeper eccentricity distribution, the projected separations represent a more biased tracer of the true values and could be considered lower limits. Such shifts in the abscissae of Figure \ref{fig3} would not explain the lack of disks around the targets with very close projected separations from $\theta^1$\,Ori C, nor do they seem likely to be large enough to erase the overall trend (although they indeed may adjust the basic shape). An intrinsic correlation between the masses of disks and their stellar hosts \citep[as found in the Taurus region by][]{andrews13} could account for the observed trend in Figure \ref{fig3} if the least massive stars are preferentially located near $\theta^1$\,Ori C. Unfortunately, the nature of the Orion proplyds makes a direct determination of their stellar masses exceedingly difficult. \citet{hillenbrand98} argued that stellar mass segregation in the ONC works in the opposite sense, with the highest mass stars ($\ge$5\,$M_{\odot}$) concentrated toward the cluster center. If that were the case, we should have identified an anti-correlation between disk mass and distance from $\theta^1$\,Ori C, which is clearly not observed. High optical depths could be responsible for such a correlation, if the disks located near $\theta^1$\,Ori C are smaller than the distant disks, and most of their emission comes from optically thick regions. However, no correlation has been observed between disk size and distance from $\theta^1$\,Ori C \citep{vicente}. Furthermore, a small ($\sim$\,50\,AU; the resolution of {\it HST}), completely optically thick disk would be detectable by our sensitive ALMA observations, with a flux density of $\sim$\,55\,mJy if viewed face-on, and an order of magnitude lower, $\sim$\,5.5\,mJy, if viewed nearly edge-on, suggesting the submillimeter wave optical depths are not responsible for the observed correlation. Instead, the evidence suggests that an externally driven disk evolution factor is likely responsible for the behavior in Figure \ref{fig3}. Tidal stripping by stellar encounters is not only too inefficient for substantial disk destruction in the ONC \citep{scally01,hollenbach00}, but, as \citet{mann09b} argued, the conditions required for disk-disk interactions to deplete disk masses \citep[e.g.,][]{olczak} also implicitly involve very high photoevaporation mass-loss rates. Overall, the data suggest that photoevaporative mass-loss driven by the ultraviolet radiation from $\theta^1$\,Ori C is the most dominant process responsible for the observed relationship. Theoretical models of disk photoevaporation indeed predict mass-loss rates that decrease with distance from the irradiation source \citep{johnstone98,storzer, richling00,scally01,matsuyama,adams04}. These models suggest that only low-mass ($\la$ few $M_{\rm jup}$) disks should exist within $\sim$0.01--0.03\,pc of $\theta^1$\,Ori C because of the strong extreme-UV (EUV) irradiation at those distances \citep{johnstone98,storzer,adams04}. At larger separations, $\sim$0.03--0.3\,pc, less energetic far-UV photons dominate the radiation field, resulting in lower mass-loss rates and thereby preserving more massive disks for up to a few Myr \citep[e.g.,][]{adams04}. This predicted behavior is consistent with the observations in the context of Figure \ref{fig3}. There is a clear lack of massive disks ($\gtrsim$\,3\,$M_{\rm jup}$) within 0.03\,pc of $\theta^1$\,Ori C where EUV irradiation dominates, whereas we find a wide range of disk masses (similar to what is found in low-mass star formation regions) at larger projected separations in the less destructive FUV-dominated regime. Accordingly, the potential to form a planetary system like our own in the EUV-dominated region of the ONC seems unlikely, given the substantially depleted disk masses there. If these nearby disks have not formed planets already, they may be out of luck unless dust grains have grown very large in these disks, to sizes not probed by submillimeter wavelength observations. Resolved, multi-wavelength observations of the Orion proplyds are required to investigate how far planet formation has already progressed in these young disks. It is interesting to note, however, that the fraction of disks with masses that exceed the nominal ``Minimum Mass Solar Nebula" model \citep[$\sim$10\,$M_{\rm jup}$;][]{weidenschilling} in the more distant FUV-dominated region of the ONC is essentially the same as that found in the low-mass star formation environment of Taurus \citep[$\sim$10\%;][]{andrews13}.\footnote{Although it is worth noting that there are still strong selection effects at play in the currently incomplete ONC disk mass census that will need to be revisited when forthcoming ALMA datasets become available.} Overall, these observations support the idea that the strength of the local EUV irradiation field has profound environmental consequences on the potential for giant planet formation in the centers of massive star-forming regions. In ALMA Cycle 1, we expect to observe the disks around 300 stars in the ONC, including 160 {\em HST}-identified proplyds. This larger scale study will allow us to survey disks across a range of distances out to 1.6\,pc from $\theta^1$\,Ori\,C, to probe different conditions in this massive star forming environment and uncover the overall disk fraction and the potential for forming planetary systems like our own.
14
3
1403.2026
We present Atacama Large Millimeter/submillimeter Array (ALMA) observations of protoplanetary disks ("proplyds") in the Orion Nebula Cluster. We imaged five individual fields at 856 μm containing 22 Hubble Space Telescope (HST)-identified proplyds and detected 21 of them. Eight of those disks were detected for the first time at submillimeter wavelengths, including the most prominent, well-known proplyd in the entire Orion Nebula, 114-426. Thermal dust emission in excess of any free-free component was measured in all but one of the detected disks, and ranged between 1 and 163 mJy, with resulting disk masses of 0.3-79 M <SUB>jup</SUB>. An additional 26 stars with no prior evidence of associated disks in HST observations were also imaged within the 5 fields, but only 2 were detected. The disk mass upper limits for the undetected targets, which include OB stars, θ<SUP>1</SUP> Ori C, and θ<SUP>1</SUP> Ori F, range from 0.1 to 0.6 M <SUB>jup</SUB>. Combining these ALMA data with previous Submillimeter Array observations, we find a lack of massive (gsim3 M <SUB>jup</SUB>) disks in the extreme-UV-dominated region of Orion, within 0.03 pc of θ<SUP>1</SUP> Ori C. At larger separations from θ<SUP>1</SUP> Ori C, in the far-UV-dominated region, there is a wide range of disk masses, similar to what is found in low-mass star forming regions. Taken together, these results suggest that a rapid dissipation of disk masses likely inhibits potential planet formation in the extreme-UV-dominated regions of OB associations, but leaves disks in the far-UV-dominated regions relatively unaffected.
false
[ "disk masses", "resulting disk masses", "disks", "associated disks", "protoplanetary disks", "θ", "jup</SUB", "OB stars", "Orion", "OB associations", "previous Submillimeter Array observations", "The disk mass upper limits", "HST)-identified proplyds", "low-mass star forming regions", "the detected disks", "HST observations", "the Orion Nebula Cluster", "submillimeter wavelengths", "potential planet formation", "Hubble Space Telescope" ]
9.945139
11.842865
-1
818378
[ "Abazajian, Kevork N." ]
2014PhRvL.112p1303A
[ "Resonantly Produced 7Â keV Sterile Neutrino Dark Matter Models and the Properties of Milky Way Satellites" ]
143
[ "Center for Cosmology, Department of Physics and Astronomy, University of California, Irvine, Irvine, California 92697, USA" ]
[ "2013arXiv1304.0759D", "2014AASP....4....9I", "2014AnP...526..294B", "2014ApJ...792...24P", "2014JCAP...05..033F", "2014JCAP...10..001A", "2014JCAP...10..013C", "2014JCAP...11..033C", "2014JCAP...12..004H", "2014JHEP...08..086G", "2014JHEP...09..006C", "2014JHEP...10..146C", "2014JHEP...11..018D", "2014JHEP...11..140F", "2014PhLB..733..217P", "2014PhLB..734...96N", "2014PhLB..734..183I", "2014PhLB..735...92C", "2014PhLB..738..380F", "2014PhRvD..89k3004S", "2014PhRvD..89k5011B", "2014PhRvD..89l1302C", "2014PhRvD..89l7305A", "2014PhRvD..90a1701B", "2014PhRvD..90a3007T", "2014PhRvD..90b3535K", "2014PhRvD..90b3540C", "2014PhRvD..90c5002D", "2014PhRvD..90c5005R", "2014PhRvD..90c5009C", "2014PhRvD..90c5014D", "2014PhRvD..90d5030R", "2014PhRvD..90l3531D", "2014PhyOJ...7..128A", "2014arXiv1405.2931D", "2014arXiv1408.1157N", "2014arXiv1409.3656P", "2015EPJC...75..471K", "2015JCAP...01..019C", "2015JCAP...02..020M", "2015JCAP...04..006V", "2015JCAP...05..053C", "2015JCAP...06..011M", "2015JCAP...07..023A", "2015JCAP...09..066P", "2015JHEP...01..006A", "2015JHEP...08..072A", "2015JHEP...10..046D", "2015JHEP...11..171G", "2015JPhCS.627a2008C", "2015MNRAS.446.2760M", "2015MNRAS.450.2143J", "2015MNRAS.451.1573L", "2015MNRAS.451.2447U", "2015MNRAS.451.3117S", "2015PhLB..742..249K", "2015PhLB..749..283M", "2015PhRvD..91d2005M", "2015PhRvD..91f3502L", "2015PhRvD..91g4005S", "2015PhRvD..91k5013P", "2015PhRvD..91l3509A", "2015PhRvD..92d3503N", "2015PhRvD..92f3502A", "2015PhRvD..92g5005B", "2015PhRvD..92i5018R", "2015PhRvD..92i6013C", "2015PhRvD..92j3509P", "2015PhRvL.115p1301B", "2015RvMP...87..531G", "2015arXiv150201320G", "2015arXiv150203399H", "2016AASP....6....3I", "2016Ap&SS.361..116C", "2016ApJ...818...90M", "2016ApJ...831...55B", "2016JCAP...01..031H", "2016JCAP...02..016D", "2016JCAP...04..003M", "2016JCAP...04..059S", "2016JCAP...06..006W", "2016JCAP...06..011L", "2016JCAP...06..017R", "2016JCAP...08..012B", "2016JCAP...11..038K", "2016JPhG...43a3001M", "2016MNRAS.455..318B", "2016MNRAS.456.4346H", "2016MNRAS.459.1489B", "2016MNRAS.461...60L", "2016MNRAS.462.4473N", "2016MNRAS.463.3848B", "2016PhRvD..94d3515V", "2016PhRvD..94h3505J", "2016PhRvL.116c1301S", "2016arXiv160505326P", "2016arXiv160508788Q", "2016arXiv160702931M", "2016arXiv160904671L", "2016arXiv160907081L", "2017ApJ...836...61M", "2017ApJ...836..217C", "2017JCAP...01..025A", "2017JCAP...06..008D", "2017JCAP...06..036G", "2017MNRAS.468.4285L", "2017MNRAS.472.4099K", "2017PhR...711....1A", "2017PhRvD..95d3524L", "2017PhRvD..95f3503G", "2017PhRvD..95l3002P", "2017PhRvD..96k5019B", "2017PhRvL.118f1101G", "2018ApJ...854....1M", "2018PhRvD..97b3020J", "2018PhRvD..98j3010R", "2018RPPh...81f6201R", "2019ASSP...56....1A", "2019ASSP...56...19B", "2019IJMPA..3430005K", "2019MNRAS.483.4080R", "2019MNRAS.483.4086B", "2019MNRAS.485.4071L", "2019PhRvD..99h3005N", "2019PhRvD.100h3536J", "2019PhRvD.100j3513A", "2019PrPNP.104....1B", "2019arXiv190310712P", "2020JCAP...06..008G", "2020JHEP...04..178B", "2020JMPh...11..168B", "2020MNRAS.495.3022P", "2020PhRvD.101j3011R", "2021PhRvL.126i1101N", "2021Univ....7..264B", "2021arXiv211214527C", "2022PrPNP.12203913D", "2023JCAP...09..047H", "2023JHEP...03..063D", "2023MNRAS.524.6345L", "2023arXiv230503032D", "2023arXiv230916942S", "2024arXiv240214333M" ]
[ "astronomy", "physics" ]
9
[ "95.35.+d", "14.60.Pq", "14.60.St", "98.65.-r", "Dark matter", "Neutrino mass and mixing", "Non-standard-model neutrinos right-handed neutrinos etc.", "Galaxy groups clusters and superclusters", "large scale structure of the Universe", "Astrophysics - Cosmology and Nongalactic Astrophysics", "Astrophysics - Astrophysics of Galaxies", "Astrophysics - High Energy Astrophysical Phenomena", "High Energy Physics - Phenomenology" ]
[ "1982Natur.299...37B", "1985NuPhB.249..361A", "1994PhRvL..72...17D", "1999PhRvL..82.2832S", "2000ApJ...538..473L", "2001ApJ...556...93B", "2001ApJ...562..593A", "2001PhRvD..64b3501A", "2002PhRvD..66b3526A", "2006PhRvD..73f3506A", "2006PhRvL..97s1303S", "2006PhRvL..97z1302B", "2007JHEP...07..030B", "2008JCAP...06..031L", "2008PhRvD..78b3524K", "2011APh....35..177A", "2011MNRAS.415L..40B", "2011PhRvD..83d3506P", "2012MNRAS.420.2318L", "2012MNRAS.422.1203B", "2013JCAP...03..014A", "2013NuPhB.866..177D", "2013PhRvD..88d3502V", "2014ApJ...789...13B", "2014MNRAS.437.2922P", "2014MNRAS.441L...6S", "2014MNRAS.442.1597S", "2014MNRAS.442.2487K", "2014PhRvD..89b5017H", "2014PhRvL.113y1301B", "2015PNAS..11212249W", "2016PhRvD..94h3002F" ]
[ "10.1103/PhysRevLett.112.161303", "10.48550/arXiv.1403.0954" ]
1403
1403.0954_arXiv.txt
14
3
1403.0954
Sterile neutrinos produced through a resonant Shi-Fuller mechanism are arguably the simplest model for a dark matter interpretation of the origin of the recent unidentified x-ray line seen toward a number of objects harboring dark matter. Here, I calculate the exact parameters required in this mechanism to produce the signal. The suppression of small-scale structure predicted by these models is consistent with Local Group and high-z galaxy count constraints. Very significantly, the parameters necessary in these models to produce the full dark matter density fulfill previously determined requirements to successfully match the Milky Way Galaxy's total satellite abundance, the satellites' radial distribution, and their mass density profile, or the "too-big-to-fail problem." I also discuss how further precision determinations of the detailed properties of the candidate sterile neutrino dark matter can probe the nature of the quark-hadron transition, which takes place during the dark matter production.
false
[ "dark matter", "the full dark matter density fulfill", "the candidate sterile neutrino dark matter", "the dark matter production", "a dark matter interpretation", "the recent unidentified x-ray line", "Sterile neutrinos", "high-z galaxy count constraints", "the Milky Way Galaxys total satellite abundance", "objects", "Local Group", "place", "their mass density profile", "the satellites radial distribution", "the simplest model", "the quark-hadron transition", "a resonant Shi-Fuller mechanism", "Shi-Fuller", "how further precision determinations", "small-scale structure" ]
7.328873
-1.594242
42
1613820
[ "ZuHone, J. A.", "Brunetti, G.", "Giacintucci, S.", "Markevitch, M." ]
2015ApJ...801..146Z
[ "Testing Secondary Models for the Origin of Radio Mini-Halos in Galaxy Clusters" ]
31
[ "Astrophysics Science Division, Laboratory for High Energy Astrophysics, Code 662, NASA/Goddard Space Flight Center, Greenbelt, MD 20771, USA", "INAF Istituto di Radioastronomia, via Gobetti 101, I-40129 Bologna, Italy", "Department of Astronomy, University of Maryland, College Park, MD, 20742-2421, USA", "Astrophysics Science Division, Laboratory for High Energy Astrophysics, Code 662, NASA/Goddard Space Flight Center, Greenbelt, MD 20771, USA" ]
[ "2014ApJ...795...73G", "2015A&A...575A...8D", "2015ApJ...811..111M", "2016A&A...589A..33A", "2016MNRAS.459.2940K", "2018MNRAS.480.5113M", "2018SSRv..214..122D", "2019ApJ...880...70G", "2019MNRAS.486L..80K", "2019SSRv..215...16V", "2020A&A...640A..37I", "2021A&A...646A..38T", "2021A&A...647A..50C", "2021A&A...649A..37B", "2021ApJ...914...73Z", "2021Galax...9...91Z", "2021MNRAS.508.2862P", "2021MNRAS.508.3995B", "2022A&A...668A.107D", "2022MNRAS.512.2157C", "2022MNRAS.512.4210R", "2022MNRAS.514.5969K", "2022MNRAS.515.1871R", "2023A&A...678A.133B", "2023MNRAS.518.4595S", "2023MNRAS.520.4410T", "2023MNRAS.520L...1P", "2023MNRAS.524.6052R", "2023Univ....9..319W", "2024ApJ...961...15N", "2024ApJ...961..133G" ]
[ "astronomy" ]
3
[ "galaxies: clusters: general", "magnetohydrodynamics: MHD", "radio continuum: galaxies", "X-rays: galaxies: clusters", "Astrophysics - Cosmology and Nongalactic Astrophysics", "Astrophysics - Astrophysics of Galaxies", "Astrophysics - High Energy Astrophysical Phenomena" ]
[ "1992ApJ...388L..49B", "1999APh....12..169B", "1999ApJ...520..529S", "2000A&A...362..151D", "2001ApJ...556L..91S", "2001ApJ...559...59M", "2001CoPhC.141...17M", "2001MNRAS.320..365B", "2002A&A...386..456G", "2004A&A...413...17P", "2005MNRAS.363.1173B", "2006ApJ...650..102A", "2006PhRvD..74c4018K", "2007A&A...470L..25G", "2007ASPC..379..221P", "2007MNRAS.378..245B", "2007PhR...443....1M", "2008A&A...486L..31C", "2008ApJ...675L...9M", "2008MNRAS.385.1211P", "2009A&A...499..371G", "2009JCoPh.228..952L", "2010A&A...513A..30B", "2010A&A...514A..76M", "2010A&A...516A..32G", "2010ApJ...717..908Z", "2010ApJ...717L..71A", "2010ApJ...719L..74K", "2010ApJ...722..737K", "2010arXiv1011.0729K", "2011A&A...525L..10G", "2011A&A...527A..99E", "2011ApJ...743...16Z", "2011ApJS..192....9T", "2011MNRAS.410..127B", "2011MNRAS.413.2057R", "2012A&A...541A..99A", "2012A&A...544A.103V", "2012ApJ...746...53F", "2012MNRAS.419.1338R", "2012MNRAS.421.3375V", "2013ApJ...762...78Z", "2013ApJ...777..163H", "2013MNRAS.434.2209W", "2013MNRAS.436L..84G", "2014A&A...567A..93P", "2014ApJ...781....9G", "2014ApJ...787...18A", "2014ApJ...795...73G", "2014IJMPD..2330007B", "2014MNRAS.438..124Z", "2014MNRAS.439.2662V", "2015ApJ...798...90Z" ]
[ "10.1088/0004-637X/801/2/146", "10.48550/arXiv.1403.6743" ]
1403
1403.6743_arXiv.txt
} A number of relaxed, cool-core clusters are hosts to faint, diffuse synchrotron radio sources called ``radio minihalos'', with a radius comparable to the size of the cooling region ($r \simlt$~100-300~kpc) and a steep spectrum ($\alpha \simgt 1$; $S_{\nu} \propto \nu^{-\alpha}$). They arise from cosmic-ray electrons (hereafter CRe) emitting in the $\sim\mu$G magnetic fields in the cluster. These are relatively rare sources, with currently only around 10 clusters with confirmed detections. Examples include Perseus \citep{bur92}, A2029 \citep{gov09}, Ophiuchus \citep{gov09,mur10}, RXC J1504.1-0248 \citep{gia11}, and RXJ 1347-1145 \citep{git07}, along with the newest detections in \citet{gia14a}. The origin of radio mini-halos in cool core clusters, as well as their possible connection with giant radio halos, is still unclear \citep[see, e.g.][for a review]{bru14}. Though the clusters which host these sources typically also have active galactic nuclei, due to the fact that the radiative timescale of the electrons at the required energies for the observed emission ($\sim{10^8}$~years) is much shorter than the time required for these electrons to diffuse across the cooling region, they alone cannot account for the origin of minihalos. Most cool-core clusters also possess spiral-shaped cold fronts \citep{ghi10}, which are believed to be the product of the central cold gas ``sloshing'' in the cluster potential minimum due to interactions with subclusters. A number of simulation works have investigated the importance of this phenomenon \citep[e.g.][]{AM06,zuh10,rod11a}. Most relevant for this work, in \citet[][hereafter ZML11]{zuh11} we showed that such magnetic field amplification occurs near the cold front surfaces due to the local velocity shear, and that the magnetic field within the core region is stronger than it was previous to the onset of sloshing. The relevance of gas sloshing to radio minihalos was first pointed out by \citet{maz08}, who discovered a correlation between sloshing cold fronts and minihalo emission in two galaxy clusters. A number of other examples have since been observed \citep[e.g.][]{hla13,gia14a}. In these cases, the minihalo emission is well-confined to the area on the sky delineated by the cold front surfaces, suggesting a causal connection between sloshing motions and radio minihalos. One possible origin for the CRe which produce the observed synchrotron radiation of minihalos is reacceleration of a lower energy population of CRe due to turbulence \citep{git02}. In a previous work \citep[][hereafter Z13]{zuh13}, we investigated the possibility that radio minihalos are due to turbulent reacceleration in sloshing cluster cores. The sloshing motions produced turbulent velocities of $\delta{v} \sim 100-200~{\rm km~s^{-1}}$ on scales of 10s of kpc within the sloshing region. We showed that the radio emission produced in these simulations was consistent with the properties of observed minihalos where sloshing cold fronts are also found. In particular, the minihalo emission was confined to the core region and possessed a steep spectrum. There is an alternative perspective on the origin diffuse radio emission in clusters of galaxies, known as the ``hadronic'' or ``secondary'' model \citep[see, e.g.][]{pfr04, kes10a, fuj12}. Cosmic-ray protons (hereafter CRp) are believed to fill the cluster volume, having been accelerated from the thermal population to relativistic energies by supernovae, AGN, and shocks associated with cosmological structure formation \citep[for a review see][]{bru14} These protons will undergo interactions with the thermal proton population, producing pions which will decay into secondary products, including cosmic-ray electrons and positions: \begin{eqnarray} p_{\rm th} + p_{\rm CR} & \rightarrow & \pi^0 + \pi^+ + \pi^- \\ \pi^\pm & \rightarrow & \mu^\pm + \nu_\mu/\bar{\nu}_\mu \rightarrow e^\pm + \nu_e/\bar{\nu}_e + \nu_\mu + \bar{\nu}_\mu \nonumber \\ \pi^0 & \rightarrow & 2\gamma \nonumber \end{eqnarray} Since the CRp have very long radiative loss times compared with the age of the cluster, these interactions should continuously provide a fresh population of CRe at a range of energies. In this model, it is these CRe that produce the observed radio emission. These hadronic interactions should also produce a flux of $\gamma$-rays, but so far no confirmed detections of such emission from galaxy clusters have been made. Upper limits on the $\gamma$-ray flux from several experiments \citep[see, e.g.][]{ale12,fer13,pro13} indicate the ratio of the CRp energy to the thermal energy of the gas is at most $\sim$1-2\%. Still, contrary to the case of nearby giant radio halos, these limits do not put significant tension on a secondary origin of mini-halos \citep[][]{bru14}. For example, \citet{zan14} proposed that mini-halos are primarily of hadronic origin, while giant radio halos experience a transition from a central hadronic emission to a leptonic emission component in the external regions due to CRe reacceleration. In the hadronic model, under relatively quiescent conditions, the electron spectrum will reach a steady-state condition where the energy losses and the gains due to injection balance each other out \citep{sar99}. If the spectrum of CRp is a power law with spectral index $\alpha_p$, at high energies, where the radiative (synchrotron and inverse-Compton) losses dominate, the CRe spectrum has the form $N(E) \propto E^{-\alpha_e}$, which results in a synchrotron emissivity that depends on the properties of the CRp and the plasma as \begin{equation}\label{eqn:ss_j} j_\nu \propto n_{\rm th}\epsilon_{\rm CRp}\frac{B^{1+\alpha}}{B^2+B_{\rm CMB}^2}\nu^{-\alpha} \end{equation} where $n_{\rm th}$ is the number density of thermal particles, $B$ is the magnetic field strength, $B_{\rm CMB} \approx 3.25(1+z)^2~\mu$G is the equivalent magnetic field strength of the CMB, $\alpha$ is the synchrotron spectral index, and $\epsilon_{\rm CRp}$ is the energy density of CRp. Under the steady-state assumption and where radiative losses dominate, $\alpha = \alpha_p/2 \sim 1$, using a canonical value $\alpha_p \sim 2-2.5.$ This implies that for $B \gg B_{\rm CMB}$ the radio emission is roughly independent of the magnetic field strength and for $B \ll B_{\rm CMB}$ it is suppressed. Though this spectral index value is consistent with the range of spectral indices found in minihalo sources, this simplified picture implies that the synchrotron spectrum should have the same slope everywhere throughout the minihalo. So far, the few available spatially-resolved observations of the minihalo spectral slopes indicate that this is likely not the case \citep{sij93,mur10,git13,gia14b}, although more detailed analyses are needed. Under stationary conditions, and if diffusion and/or transport of CRp and CRe is not important, this power-law slope should extend to high frequencies, without any steepening (distinct from the turbulent acceleration model, which predicts steepening at high frequencies due to the balance between reacceleration and losses on the CRe). However, the conditions of cluster cores with minihalos probably do not lend themselves to a such a steady-state assumption, given the existence of cold fronts which betray the presence of sloshing motions. As previously noted (ZML11), sloshing motions can significantly amplify the magnetic field strength on short timescales. \citet{kes10b} demonstrated that a rapidly increasing magnetic field strength will result in a steeper spectrum than the steady-state case (so long as $B \gg B_{\rm CMB}$). In the present case, magnetic field amplification near cold front surfaces would cause spectral steepening of the radio emission moving outward along the spiral arm of the sloshing cold fronts, away from the cluster center. The same magnetic field amplification has possible implications for the extent of the radio emission as well. \citet{kes10a} predicted that as a result of the strengthening of the magnetic field, the radio emission would have a steep gradient at the front surface, since inside the front $B \gg B_{\rm CMB}$, and outside, $B \ll B_{\rm CMB}$. The appeal of this model for the steepening of the radio spectrum and the radial extent of minihalos is that it relies on a process, namely the amplification of the magnetic field by sloshing, that is strongly evidenced by previous theoretical and simulation work, and is independent of the details of the underlying CRp spectrum and spatial distribution. For this reason, we believe it deserves close examination. In this work, we use the same cluster setup as in Z13 to determine if the amplification of the magnetic field strength by the sloshing motions is sufficient {\it by itself} to produce the two distinct characteristics mentioned above, e.g., emission confined to the cold front region with a steep spectrum, steeper than that expected from the canonical CRp slope under stationary conditions. We follow the evolution of CRe spectra associated with passive tracer particles which are advected with the flow of gas, determining the injection into and radiative losses on the spectrum along each trajectory. In this way we can determine how deviations from a steady-state configuration affect the surface brightness of the minihalo emission as well as its spectral shape. This paper is organized as follows. In Section 2 we describe our method for evolving the CRe spectra and our assumptions regarding their injection by hadronic processes. In Section 3 we detail the results of our simulations. In Section 4 we discuss the implications of this work, and in Section 5 we make our conclusions. Thoughout this work we assume a flat $\Lambda$CDM cosmology with $h$ = 0.7 and $\Omega_{\rm m}$ = 0.3.
\subsection{The Radial Extent of Minihalos}\label{sec:radial_extent} In some observed minihalos, the radio emission is strongly confined to the core region, or the envelope of the cold fronts, if they are clearly observed. In our simulations, we find that the radial profile of the minihalo emission typically decreases at a more rapid rate approaching the cold fronts, but this is not always the case (in particular the NW profile from Figure \ref{fig:profiles}). \begin{figure} \begin{center} \plotone{compare_spectra.eps} \caption{The total radio spectrum within a radius of 300~kpc from the cluster center for the hadronic model from this work and the reacceleration model from Z13 at several epochs.\label{fig:compare_spectra}} \end{center} \end{figure} A drop in radio emission across a cold front in the secondary model primarily results from two effects: the first is the density jump across the cold front, which is typically a factor of $\sim$1.5-2 for most cold fronts. The contribution of the density jump to the radio drop will be modest, since for the case where the CRp energy density is proportional to the thermal density, $j_\nu \propto n_{\rm th}P_{\rm th}$, and the thermal pressure is roughly continuous across cold front surfaces (except in a very thin layer at the cold front surfaces where the enhanced magnetic pressure due to shear amplification may become significant). The main contribution to the drop in radio emission will arise from the difference in the magnetic field across the front \citep{kes10a}. Below the front, the field will be amplified to $B > B_{\rm CMB}$, and the radio emission is roughly independent of the magnetic field strength. When the fronts reach a large radius, above the front the unamplified field will likely be weak due to the decline of magnetic field strength with radius, with $B < B_{\rm CMB}$, and the radio emission will drop off rapidly as the square of the field. Figure \ref{fig:scaledB} shows the mass-weighted projected magnetic field at the epoch $t = 3.5$~Gyr scaled by $B_{\rm CMB}$. Throughout the volume of the sloshing region, the magnetic field has been amplified, and is stronger than $B_{\rm CMB}$. For the most part, just outside of the cold fronts the field drops to $B < B_{\rm CMB}$. However, the region to the northwest is filled with strong magnetic field outside of the cold fronts, due to strong velocity shears in that region. Along projected radii in that direction, there will be no drop in emission as one crosses the cold front surface. Clearly, in the hadronic model the radial extent of the minihalo emission depends critically on the magnetic field structure. \begin{figure*} \begin{center} \includegraphics[width=0.45\linewidth]{spec_index_map_0050.eps} \includegraphics[width=0.45\linewidth]{spec_index_map_0100.eps} \includegraphics[width=0.45\linewidth]{spec_index_map_0150.eps} \includegraphics[width=0.45\linewidth]{spec_index_map_0200.eps} \caption{Spectral index map at several epochs (between the frequencies of 327 and 1420~MHz) with 327~MHz radio contours overlaid (contours are taken from Figure \ref{fig:327_maps}). Each panel is 750~kpc on a side.\label{fig:spec_index_maps}} \end{center} \end{figure*} \begin{figure*} \begin{center} \includegraphics[width=0.45\linewidth]{spec_index_map_low_0050.eps} \includegraphics[width=0.45\linewidth]{spec_index_map_low_0100.eps} \includegraphics[width=0.45\linewidth]{spec_index_map_low_0150.eps} \includegraphics[width=0.45\linewidth]{spec_index_map_low_0200.eps} \caption{Spectral index map at several epochs (between the frequencies of 60 and 153~MHz) with 327~MHz radio contours overlaid (contours are taken from Figure \ref{fig:327_maps}). Each panel is 750~kpc on a side.\label{fig:spec_index_maps_low}} \end{center} \end{figure*} This, of course, depends to a certain extent on our initial choice for the magnetic field strength. In some clusters, the magnetic field strength may be even higher than in our setup (given by $\beta \sim 100$, see Section \ref{sec:MHD}) due to turbulent-driven amplification caused by AGN, minor mergers, and accretion. More simulations with different magnetic field configurations would be needed in order to determine the conditions under which minihalos would have steep drops in emission at cold front surfaces in the hadronic model. The significance of the drop in emission due to a sharp decrease in the magnetic field strength across the front is also affected by the strong redshift dependence of the CMB energy density, with $B_{\rm CMB} \propto (1+z)^2$. \subsection{Total Spectrum and Spectral Steepening Along the Cold Fronts} According to \citet{kes10b}, rapid magnetic field amplification can cause significant departures from steady-state conditions, which produces synchrotron spectral steepening under the hypothesis of a secondary origin of the emitting particles. If strong, these effects may have important consequences for interpreting the origin of minihalos and giant halos. In our simulations we do not see significant effects of this nature. We find that the total spectrum of the minihalo is fairly consistent with a simple power-law in agreement with the expectation based on steady-state conditions. Though we do see some steepening of the synchrotron spectrum near the cold front surfaces that we can confirm is indeed due to magnetic field amplification in these regions, we find that the effect is barely observable, even at very low observing frequencies where it should be maximized. Section 4.3.3 of \citet{kes10b} determines the change in spectral index $\alpha$ for a sudden increase in the cooling rate over the injection rate, parameterized by \begin{equation} \R = \frac{\psi_f/\Q_f}{\psi_i/\Q_i} = \frac{B_f^2+B_{\rm CMB,f}^2}{B_i^2+B_{\rm CMB,i}^2}\frac{n_{\rm th,i}C_{\rm p,i}}{n_{\rm th,f}C_{\rm p,f}}. \end{equation} From Figure 16 of \citet{kes10b}, we may determine the approximate increase in the magnetic field strength associated with a given spectral steepening, which will serve as an independent check on our model for the steepening of the CRe spectra. It is instructive to examine the necessary increase in the cooling rate to produce the steepest spectra of any of the particles in the simulation ($\alpha \sim 2$) and the steepest spectra we measure in our spectral index maps ($\alpha \sim 1.3$). For our injection index of $\alpha_p \approx 2$, a sudden increase of $\R \sim 15$ is required to steepen the synchrotron spectrum to $\alpha \sim 2$, whereas for steepening to $\alpha \sim 1.3$ a more modest increase of $\R \sim 3$ is required. What qualifies as a ``sudden'' increase? Formally, since in steady-state conditions the characteristic timescale is \citep{kes10a}: \begin{equation} t_{\rm cool} \simeq 0.13\left[\frac{4\left(\frac{B\sqrt{3}}{B_{\rm CMB,0}}\right)^{-3/2}}{1+\left(\frac{B}{B_{\rm CMB,0}}\right)^{-2}}\right]\nu_{1.4}^{-\frac{1}{2}}(1+z)^{-\frac{7}{2}}~{\rm Gyr}, \end{equation} any sudden increase in the cooling rate must occur on a timescale shorter than this in order to significantly affect the spectrum. For $B \simlt B_{\rm CMB}$, $t_{\rm cool} \sim 0.1-0.2$~Gyr for electrons emitting at frequencies between 327 and 1400~MHz. Motivated by this, we examine the trajectories of a few tracer particles with extreme ($\alpha \sim 2$) spectral indices, to see if these conditions prevail. Figure \ref{fig:cooling_evolution} shows the evolution in $\R$ for three tracer particles with spectral indices $\alpha \sim 2$. $\R$ for each particle is normalized to unity at time $t_i$, and all times are given with respect to the epoch $t_f$ at which the spectral index is measured, marked by the black dashed line. Each of these particles indeed experiences a sudden increase in $\R$, followed by a rapid decrease. The dot-dashed lines in the figure indicate the necessary $\R$ to steepen the spectrum to the given spectral index from a steady-state, estimated from Figure 16 in \citet{kes10b}. All of the curves show an increase in $\R$ that is at least this much, and this occurs over roughly a cooling time ($t \sim 0.13$~Gyr), with the most significant increase occurring over a small fraction of that time. We therefore find that the evolution of the CRe energies for the steepest-spectrum tracer particles in our simulation is consistent with the results of \citep{kes10b}. The fact that we find so few tracer particles with very steep spectra indicates that these conditions (required to produce such steepening) do not prevail throughout most of the cluster core. Though there is significant magnetic field amplification (see Figure \ref{fig:scaledB}), it does not usually occur fast enough to steepen the spectrum significantly. \subsection{Limitations of this Work}\label{sec:limitations} In this work, we have adopted a simple model for hadronically originating CRe, by relaxing the typical steady-state assumption and allowing for the time evolution of the CRe spectrum due to changes in the local gas and CRp density and the magnetic field strength, as well as the evolution of the CMB energy density with redshift. Relaxing the steady-state assumption is meaningful in our simulation, due to our high spatial and time resolution which are able to resolve adequately the amplification of the magnetic field. As a significant simplification, we also assumed that at any given epoch the energy density of CRp was proportional to the energy density of the thermal gas, and that the input spectrum of the CRp is a constant power-law spectrum (excepting the cutoff at the low-energy threshold). This approach is therefore limited, as it relies on a simplified treatment of the spatial distribution and energetics of the CRp. First, we are not able to model in a self-consistent fashion the transport and diffusion of CRp and its connection with the gas dynamics in the sloshing region. Turbulence might transport CRp on a scale that could be of the order of the cluster core, potentially inducing a spatial distribution of CRp that is broader than that of the thermal plasma. Momentum-diffusion of CRs that is mediated by the scattering of particles with MHD waves in the plasma is a different process that is more difficult to model due to the uncertainties in the physics of (small-scale) turbulence in these environments. In addition, \citet{ens11} and \citet{wie13} proposed that super-Alfv\'enic CR streaming is possible in the ICM. In particular, \citet{wie13} calculated the suppression of the streaming instability in a turbulent flow showing that under these conditions the self-generated waves do not limit the particle drift velocity to the Alfv\'en speed. However, this does not automatically imply that streaming is efficient, because the background turbulence (necessary to suppress the instability) still provides a source of scattering that may make the transport of CRs diffusive and potentially inefficient \citep[see][for a discussion]{bru14}. The combination of these effects may significantly affect the expected properties of minihalos with respect to our simplified treatment. Turbulent transport and streaming would be the most important effects, and they typically produce a broader spatial distribution of CRp with respect to our simplified model. However, if these effects are important, they would produce synchrotron distribution even broader than those obtained by our model, making the difference between reacceleration and hadronic models even more relevant. In particular, the radial drops in hadronic models could become even smoother, increasing the tension with observations. Another possibility that we cannot rule out, which would produce the opposite effect, is that the CRp which are ultimately responsible for the radio emission are generated in the cluster center (e.g., from AGN activity), and then diffuse outward, but are confined to the region bounded by the cold fronts due to the fact that the field direction is largely tangential at the front surfaces (ZML11) and thus would prevent diffusion across them. To investigate these and other effects, future studies of sloshing and minihalos based on simulation methods that self-consistently follow the transport of CRp will be needed. Second, given the simplified assumptions in our model with respect to the CRp, the spatial and spectral distribution of the radio emission that is produced in our model is determined almost entirely by the spatial distribution and time evolution of the magnetic field. It is therefore crucial that we adequately resolve this evolution and are not missing any significant sources of further amplification. The magnetic field in our cluster core will be amplified by shearing motions localized at the cold fronts, as well as turbulence generated by sloshing. \citet{kes10} identified shear flows at cold front surfaces as the most relevant mechanism for the amplification of the magnetic field, which was confirmed in ideal MHD simulations by ZML11. However, turbulence will also play a role. In our inviscid simulation, turbulent dissipation is entirely numerical, and determined by the spatial resolution. Therefore, though we resolve gas motions down to scales of $\sim$kpc, the scale at which dissipation of the turbulent cascade begins to set in is slightly larger, around $\sim$10~kpc (see Z13 and references therein). In any case, any further field amplification provided by turbulent motions on scales that are not resolved by our simulations is likely to be a small fraction of what we are able to accomplish with our high-resolution simulation. The turbulence is injected via the sloshing motions on scales of $\sim$100~kpc, and cascades down to the previously mentioned scales where it is dissipated by our finite resolution. The energy flux through the turbulent cascade, $\epsilon_t \propto v_\ell^3/\ell$ (assuming a Kolmogorov spectrum) is constant with respect to the scale $\ell$. Given this fact, the kinetic energy available for conversion into magnetic energy at these unresolved scales will be small compared with the energy available at the scales we do resolve. Therefore, we do not expect that there is any significant room for further amplification by turbulence by resolving these smaller scales. Even an increase in the magnetic field energy by a few tens of percent would not significantly change our conclusions. Additionally, our previous work (Z13) showed that the kinetic energy in turbulence generated by sloshing on scales less than $\sim$30~kpc is estimated to be at most a few percent of the thermal energy, and even if all of this energy were to be transferred to the magnetic field it would still not be as large as the contribution from the larger-scale sloshing motions (see Figure 24, Z13). In the Appendix, we use a resolution study to demonstrate that we have converged with respect to the overall properties of the magnetic field within the core region. However, it is the case that our specific initial value for the overall ratio of the thermal to the magnetic pressure, $\beta$, is a free parameter in our simulations. ZML11 investigated a range of initial values for $\beta$ (over an order of magnitude from $\sim$100-6400), and demonstrated that the average magnetic field strength within the core region was similar at late times, regardless of the initial field strength. We may predict that for initial $\beta$ higher than our value of 100 (a weaker initial field), that the contrast between the field strength inside and outside the cold front would be larger, and there would be a sharper decrease in the radio emission at the cold front surface. On the other hand, for a lower initial $\beta$ (a stronger initial field), the contrast between the magnetic field strength inside and outside the front would be smaller, and the drop in radio emission would be less significant. Third, we assumed a population of CRp with a single power-law spectrum. This implies an injection spectrum of CRe that is power-law in form, and that any steepening of the radio spectrum at higher frequencies comes from strong departures from the steady-state condition. Instead, steepening of the CRe spectrum may reflect possible steepening of the CRp spectrum. Although the momentum distribution of CRs is expected to be a power law in a very broad range of astrophysical situations, this is certainly a simplification. In our conditions, momentum-dependent spatial transport of CRs could change the spectrum of CRp. For example, \citet{wie13} discussed the possibility of a steepening of the CRp spectral distribution resulting from an efficient momentum-dependent diffusion. This is also not included in our simulations. However, in non steady-state conditions, this process would eventually result in a spectral flattening of the synchrotron emission with distance, because higher energy CRp would not be diffusively trapped and propagate faster faster (this effect is unavoidable in the case that we measure variations of the radio spectrum with distance). The observational consequence would be that minihalos would be typically broader if observed at higher frequencies. This differs from expectations based on other models, allowing to discriminate between different origin scenarios for the CRe. For example, reacceleration models eventually predict a spectral steepening with radius or more complex spectral variations \citep[e.g.][and references therein]{bru14}. Another limitation is that we do not include the effect of turbulent reacceleration on the spectrum of both CRp and CRe, a process that is believed to be important, for example, in the case of giant radio halos.
14
3
1403.6743
We present an MHD simulation of the emergence of a radio minihalo in a galaxy cluster core in a “secondary” model, where the source of the synchrotron-emitting electrons is hadronic interactions between cosmic-ray protons with the thermal intracluster gas, an alternative to the “reacceleration model” where the cosmic ray electrons are reaccelerated by turbulence induced by core sloshing, which we discussed in an earlier work. We follow the evolution of cosmic-ray electron spectra and their radio emission using passive tracer particles, taking into account the time-dependent injection of electrons from hadronic interactions and their energy losses. We find that secondary electrons in a sloshing cluster core can generate diffuse synchrotron emission with luminosity and extent similar to observed radio minihalos. However, we also find important differences with our previous work. We find that the drop in radio emission at cold fronts is less prominent than that in our reacceleration-based simulations, indicating that in this flavor of the secondary model the emission is more spatially extended than in some observed minihalos. We also explore the effect of rapid changes in the magnetic field on the radio spectrum. While the resulting spectra in some regions are steeper than expected from stationary conditions, the change is marginal, with differences in the synchrotron spectral index of {Δ }α ≲ 0.15-0.25, depending on the frequency band. This is a much narrower range than claimed in the best-observed minihalos and produced in the reacceleration model. Our results provide important suggestions to constrain these models with future observations.
false
[ "observed radio minihalos", "radio emission", "diffuse synchrotron emission", "secondary electrons", "core sloshing", "hadronic interactions", "electrons", "passive tracer particles", "future observations", "the cosmic ray electrons", "the reacceleration model", "a radio minihalo", "cosmic-ray electron spectra", "the secondary model", "their radio emission", "a galaxy cluster core", "cold fronts", "cosmic-ray protons", "a sloshing cluster core", "some observed minihalos" ]
14.495003
4.767853
152
483007
[ "Zhang, Shaohua", "Wang, Huiyuan", "Wang, Tinggui", "Xing, Feijun", "Zhang, Kai", "Zhou, Hongyan", "Jiang, Peng" ]
2014ApJ...786...42Z
[ "Outflow and Hot Dust Emission in Broad Absorption Line Quasars" ]
32
[ "Polar Research Institute of China, 451 Jinqiao Road, Shanghai 200136, China", "Key Laboratory for Research in Galaxies and Cosmology, University of Science and Technology of China, Chinese Academy of Sciences, Hefei, Anhui 230026, China", "Key Laboratory for Research in Galaxies and Cosmology, University of Science and Technology of China, Chinese Academy of Sciences, Hefei, Anhui 230026, China", "Key Laboratory for Research in Galaxies and Cosmology, University of Science and Technology of China, Chinese Academy of Sciences, Hefei, Anhui 230026, China", "Key Laboratory for Research in Galaxies and Cosmology, Shanghai Astronomical Observatory, Chinese Academy of Sciences, 80 Nandan Road, Shanghai 200030, China", "Polar Research Institute of China, 451 Jinqiao Road, Shanghai 200136, China; Key Laboratory for Research in Galaxies and Cosmology, University of Science and Technology of China, Chinese Academy of Sciences, Hefei, Anhui 230026, China", "Key Laboratory for Research in Galaxies and Cosmology, University of Science and Technology of China, Chinese Academy of Sciences, Hefei, Anhui 230026, China" ]
[ "2014MNRAS.443.2532H", "2015ApJ...802...92Z", "2015ApJ...814..150W", "2015ApJ...815..113Z", "2015ApJS..217...11L", "2016ApJ...819...99S", "2016ApJ...824..106W", "2016ApJ...829...96S", "2016MNRAS.455..136V", "2016MNRAS.459.2472R", "2016MNRAS.461..647C", "2016arXiv161103733G", "2017ApJ...836...86Z", "2018AJ....156....4Z", "2018ApJ...863..198P", "2019ApJ...879..123L", "2019MNRAS.487.3884C", "2019Natur.573...83Z", "2021ApJ...916...86L", "2021ApJS..255...12Y", "2021MNRAS.501.3061T", "2021MNRAS.504.3187M", "2021MNRAS.508.1305Y", "2021RAA....21...65L", "2022A&A...663A..63Z", "2022ApJ...926...60V", "2022ApJ...930....5Y", "2022SciA....8.3291H", "2023A&A...671A..34S", "2023ApJ...952...44B", "2024ApJ...966...87L", "2024ApJS..271...61Y" ]
[ "astronomy" ]
14
[ "dust", "extinction", "galaxies: nuclei", "infrared: galaxies", "quasars: absorption lines", "quasars: general", "Astrophysics - Astrophysics of Galaxies", "Astrophysics - Cosmology and Nongalactic Astrophysics" ]
[ "1982ApJ...263...79G", "1987ApJ...320..537B", "1987ApJS...63..615N", "1988ApJ...325...74S", "1988MNRAS.234..297S", "1991ApJ...373...23W", "1992ApJ...385..460E", "1992ApJ...390...39S", "1993ApJ...411..594N", "1993ApJ...413...95V", "1993ApJS...88..357K", "1994ApJS...95....1E", "1995ApJ...447..512K", "1996ApJ...473..781N", "1998ApJ...500..798H", "1999A&A...344..433C", "1999ARA&A..37..487H", "2000AJ....120.1579Y", "2000ApJ...545...63E", "2001AJ....122..549V", "2001ApJ...546..775B", "2002ApJ...565...78B", "2002ApJ...567L.107E", "2002ApJ...569..641L", "2002ApJ...578L..31T", "2002ApJS..141..267H", "2003AJ....125.1784H", "2003AJ....126.2594R", "2004ApJ...611..107L", "2004ApJ...611..125L", "2004MNRAS.352.1390M", "2005ARA&A..43..769V", "2005ApJ...625...95C", "2006A&A...459...85N", "2006ApJ...639...46S", "2006ApJ...648L.101E", "2006ApJS..165....1T", "2006ApJS..166....1H", "2007A&A...476..713K", "2007ApJ...663..320F", "2007ApJ...665..990G", "2007ApJS..171....1M", "2008ApJ...672..108D", "2009ASPC..411..251M", "2009ApJ...690.1006F", "2009ApJ...692..758G", "2009ApJ...702..851W", "2009ApJ...703L...1D", "2009ApJ...705..298M", "2009ApJ...707.1334W", "2009MNRAS.399.2231S", "2010AJ....139.2360S", "2010AJ....140.1868W", "2010ApJ...714..367Z", "2010MNRAS.401....7H", "2010MNRAS.405.1303A", "2011AJ....141..167R", "2011ApJ...737L..36M", "2011ApJ...738...85W", "2011ApJS..194...45S", "2011MNRAS.410..860A", "2012ApJ...751L..23W", "2012MNRAS.427.1135D", "2013AJ....145..157J", "2013AJ....146...78G", "2013ApJ...763L..18W", "2013ApJ...768...22Z", "2013ApJ...771...87H", "2013ApJ...775L..15R", "2013ApJ...776L..15W", "2013MNRAS.429.1494R", "2013MNRAS.432.1525B", "2013MNRAS.433.1778R", "2014MNRAS.438.1288H" ]
[ "10.1088/0004-637X/786/1/42", "10.48550/arXiv.1403.3166" ]
1403
1403.3166_arXiv.txt
Outflows appear to be a common phenomenon in quasars. Strong outflows may carry away huge amounts of material, energy and angular momentum and are believed to be one of the most important feedback processes connecting their central engines and host galaxies (e.g. Hopkins \& Hernquist 2006; Hopinks \& Elvis 2010). For instance, outflows are thought to be able to regulate the growth of supermassive black holes (SMBH) and star formation in host galaxies (see Antonuccio-Delogu \& Silk 2010 for a recent review), chemically enrich the interstellar medium (ISM) of host galaxies and the surrounding intergalactic medium (IGM) (e.g., Collin \& Zahn 1999; Veilleux et al. 2005). Lots of studies tried to closely tie up outflows with the quasar fundamental parameters, such as Eddington ratio (Boroson 2002; Gangly et al. 2007; Zhang et al. 2010; Wang et al. 2011; Marziani \& Sulentic 2012), the intrinsic spectral energy distribution (SED, e.g. Leighly \& Moore 2004; Fan et al. 2009; Richards et al. 2011; Baskin et al. 2013) and even gas metallicity (e.g. Wang et al. 2012). Therefore outflows are important for our understanding of the evolution of the central SMBHs and the connection with host galaxies. Quasar outflows manifest themselves in blueshifted absorption and emission, such as broad absorption lines (BALs; Weymann et al. 1991) and blueshifted broad emission lines (BELs, Gaskell 1982). It is well known that BAL quasars in general have redder ultraviolet (UV) continua than non-BAL quasars, and quasars with BALs from low-ionization state species (e.g. \mgii\ and \feii\; LoBALs) are even redder than quasars with absorption in high-ionization species (e.g., \civ\ and \nv; HiBALs) on average (Weymann et al. 1991; Brotherton et al. 2001; Reichard et al. 2003). This phenomenon is readily interpreted as a result of dust extinction (e.g. Sprayberry \& Foltz 1992; Voit et al. 1993; Reichard et al. 2003; Hewett \& Foltz 2003; Dai et al. 2008; Jiang et al. 2013). The BAL outflow launching region is suggested to be co-spatial with or outside of the BEL regions, based on the observational fact that BAL features usually obscure BELs. Dust can survive at the outer boundary of BEL regions according to the reverberation mapping results for local active galactic nuclei (Suganuma et al. 2006). It implies the possibility of dusty outflows. Recently, Grupe et al. (2013) found an interesting anti-correlation between luminosity and UV continuum slope in a variable BAL Seyfert 1, WPVS 007. It also favors a dusty outflow component moving transversely across our line of sight. The composite spectra studies suggested an SMC-like reddening law for dust associated with quasar outflows (Reichard et al. 2003; Zhang et al. 2010), though it may not always be the case (see the studies of some individual objects, e.g. Hall et al. 2002; Jiang et al. 2013). The average amount of reddening in HiBAL and LoBAL quasars are $E(B-V)\sim$0.023 and $E(B-V)\sim$0.077, respectively (Reichard et al. 2003; Zhang et al. 2010). Such large amount of radiation absorbed/scattered by dust grains must be partly re-radiated in infrared band, yielding a prominent and detectable feature. More recently, Wang et al. (2013, hereafter Paper I) found an unexpected correlation between the amount of hot dust emission relative to accretion disk emission and the blueshift of \civ\ BEL in $z\sim2$ non-BAL quasars. The correlation dramatically strengthens with increasing Eddington ratio, since outflows tend to be dominant in high Eddington ratio quasars. It strongly implies an important role of dust in the outflow physics, such acceleration mechanisms or interaction with nearby medium (Paper I), and even in the galaxy physics and cosmology (see Elvis et al. 2002). BALs and blueshifted BELs are very likely to be different appearances of the same outflow component viewed from different inclination angles (see e.g. Wang et al. 2011). The normal BELs emitted by virialized (or rotational) gases may hamper the proper measurement of outflow properties from BELs. Fortunately this effect is negligible in BAL measurements. It is thus necessary to revisit and extend the relationship between outflow properties and hot dust emission using a BAL quasar sample. In this paper, we construct a large $z\sim2$ BAL quasar sample from the Wide-field Infrared Survey Explorer (WISE; Wright et al. 2010) and the Sloan Digital Sky Survey (SDSS; York et al.2000). Following Paper I, we adopt the rest-frame NIR slope, measured from WISE data, as an indicator of the relative amount of hot dust emission. We use the BAL parameters measured from SDSS spectra, such as velocities and absorption strength, to indicate the outflow strength as commonly adopted in the literature. This paper is organized as follows. The sample construction and outflow parameter measurement are shown in Section 2. We analyze the data and present the correlations in Section 3. Finally, we discuss the possible underlying physics in Section 4, and summarize the results in Section 5. Throughout this paper, we adopt the CDM `concordance' cosmology with H0 = 70 km s$^{-1}$Mpc$^{-1}$, $\Omega_{\rm m} = 0.3$, and $\Omega_{\Lambda} = 0.7$.
In Paper I, we adopted the blueshift and asymmetry index ($BAI$) of \ion{C}{4} BELs to indicate the outflow property and velocity, and found an interesting correlation between BAI and $\beta_{\rm NIR}$. In this paper, the blueshifted velocities and absorption strength of BALs instead of BAI are used to characterize the outflow strength. Again, we find significant correlations with $\beta_{\rm NIR}$, well consistent with the results shown in Paper I. Using outflow parameters measured from BELs (e.g. BAI) and BALs (e.g. $V_{\rm max}$, $V_{\rm ave}$ and $BI$) both have advantages and shortcomings. On one hand, BELs are the integral of emission over entire volume of the outflow and represents the overall properties, while BAL troughs only hold the information of outflow along the LOS and thus are sensitive to clumpy structures. It may be the reason for that $\beta_{\rm NIR}$ is more strongly correlated with BEL parameters than BAL parameters. On the other hand, the physical meanings of BAL parameters are more straightforward and specific than those of BEL parameters. The emission from virialized gas also contributes to BELs so that it is generally hard to reliably extract outflow parameters from BELs. In contrast, the contamination in BALs from other line features is usually unimportant and easily to handle (see Section 2). It is one of the reasons that we think it is essential to use BAL quasars to revisit the relationship between outflow and hot dust. Overall, these two works together present robust evidences for the important correlation and connection between outflow strength and hot dust emission. In order to understand the origin of these correlations, we present detailed discussions in the following. In Section \ref{sec_df}, we examine whether or not a dust-free outflow, together with a hydrostatic torus model, can properly accommodate our findings. Our analysis based on current observational results suggest that this scenario is likely unviable. In Section \ref{sec_do}, we discuss the dusty outflow scenarios, which have ever been proposed in Paper I. \subsection{Dust-free outflow scenario}\label{sec_df} Suppose that hot dust emission in these high-redshift quasars has nothing to do with outflows, and is predominantly emitted by the innermost part of a hydrostatic and optically thick torus, as suggested by lots of previous studies (e.g., Neugebauer et al. 1987; Barvainis 1987; Suganuma et al. 2006; Kishimoto et al. 2007; Mor et al. 2009; Mor \& Trakhtenbrot 2011). It would be interesting to see whether or not the correlations shown in this paper and Paper I are induced by a third factor that simultaneously governs or relates to outflows and dust emission. We first discuss the inclination effect. The NIR slope, $\beta_{\rm NIR}$, measures the amount of hot dust emission relative to accretion disk emission. The disk emission is theoretically predicted to scale with the cosine of the disk inclination angle (IA). If hot dust emission is (approximately) isotropic, $\beta_{\rm NIR}$ is expected to increase with increasing IA. However, dust torus may block the hot dust emission at a large IA (Roseboom et al. 2013). It complicates the situation and makes the dependence of $\beta_{\rm NIR}$ on IA uncertain. More recently, Runnoe et al. (2013) used a sample of radio-loud quasars to investigate the inclination dependence of quasar spectral energy distribution. The composites in their figure 6 clearly present a weak dependence of $\beta_{\rm NIR}$ on IA, in the sense that a edge-on quasar spectrum tend to show a slightly larger $\beta_{\rm NIR}$ than a face-on one. We thus still assume that $\beta_{\rm NIR}$ increases with increasing IA Since the outflow covering factor is usually small, about 10\%$\sim$20\% (e.g., Tolea et al. 2002; Hewett \& Foltz 2003; Reichard et al. 2003; Trump et al. 2006), it is unlikely that outflow properties and $\beta_{\rm NIR}$ systematically and significantly vary with IA within the opening angle of outflows. However, one can still elaborately design a outflow model to match the observation. For example, the mean outflow IAs vary among quasars and the small-IA outflows are, on average, weaker than the large-IA outflows. The problem for this hypothesis is that there is no any observational evidence for it and the inclination dependence of $\beta_{\rm NIR}$ is actually weak (Runnoe et al. 2013). In addition, it also fails to account for the correlation with BAI, the outflow parameter measured from BELs (see Paper I). The inclination effect might be the other factor for the large scatter in the observed correlations. Lots of studies have found the correlations of outflow properties with quasar fundamental parameters e.g., Reichard et al. 2003; Gangly et al. 2007; Fan et al. 2009; A11; Baskin et al. 2013). Outflow strength is significantly dependent on Eddington ratio, broad band SED (e.g. ionization SED, and UV continuum slope, $\beta_{\rm UV}$) and luminosity. If these factors also have an impact on hot dust emission, one might find a piece of evidence to support the dust-free outflow scenario. As shown in this paper and previous studies (e.g. Mor \& Trakhtenbrot 2011; Paper I), both $\beta_{\rm NIR}$ and the infrared to UV flux ratio are almost independent of Eddington ratio. And the correlations between luminosity and the two hot dust indicators are either absent or negative. In addition, the correlation of outflow strength with $\beta_{\rm NIR}$ is stronger than that with $\beta_{\rm UV}$ and the correlation between the two slopes are apparently weaker (Table \ref{tab1}). Hence none of them is the factor that we are seeking. The role of ionization SED is unclear, since we do not know its relation with $\beta_{\rm NIR}$. Recently, Baskin et al. (2013) detected an interesting dependence of outflow strength on EW(\heii) and ascribed it to the impact of ionization SED. We thus check the relationship between EW(\heii) and $\beta_{\rm NIR}$. Again, no correlation is found. Another interesting factor is metallicity. Quasars harboring strong outflows tend to have high gas metallicity (Hamann 1998; Leighly 2004; Wang et al. 2009b). Recently, Wang et al. (2012) found that \ion{C}{4} blueshift increases with gas metallicity. It means that outflows are stronger in higher metallicity environment. Since dust forms more easily in higher metallicity gas, one may expect that metallicity is an appropriate factor which simultaneously influences outflow and hot dust emission. However, such naive expectation may not be true. For a typical torus, the relative amount of dust emission is determined by dust covering factor, but not dust amount. Theoretically, it is unclear how to construct a relationship between dust covering factor and metallicity. And observationally, there is no evidence supporting it. On the contrary, the required positive correlation seems inconsistent with the current observational facts that metallicity strongly increases with increasing luminosity(e.g. Hamann \& Ferland 1999; Nagao et al. 2006) and dust covering factor (indicated by the infrared to UV flux ratio) is weakly anti-correlated with luminosity (e.g. Mor \& Trakhtenbrot 2011). To sum up, the current observational results listed in this paper and the literature do not favor the dust-free outflow scenario. However further works, especially on the ionization SED and metallicity, are still required to examine our point. \subsection{Dusty outflow scenarios}\label{sec_do} In Paper I, we attempted to propose two plausible mechanisms to interpret the correlation of outflow strength with hot dust emission. Here we present further discussion. In the first mechanism, dust is intrinsic to outflows. This idea is supported by the similar locations of BAL outflows and hot dust (Elvis 2000 and references therein). There is little observational constraint on the (relative) location of BAL outflows. Since the absorption material almost always absorbs the BELs, it is usually believed that BAL outflows are co-spatial with or outside of BEL regions (BLRs). The reverberation mapping results of local AGNs suggested that hot dust is also located at the boundary of BLRs (e.g., Suganuma et al. 2006). These two implies the possibility of originally dusty outflows. Outflows may emerge from the outer region of accretion disk or even the innermost region of torus, in which the gas clouds are dusty and relatively cold. These dusty clouds are uplifted above the disk, and exposed to the central engine. The low density part is highly ionized and responsible for the blueshifted absorption and emission lines. Dust survives in the dense region and radiates in NIR band. In addition, dust perhaps forms in dense clouds embedded in outflows (Elvis et al. 2002). A strong outflow carries large amount of dust, and thus enhances the NIR emission. At the same time, the acceleration due to dust absorption and scattering make the outflow stronger, which further enhances the correlation. It is worth noting that Elitzur \& Shlosman (2006; see also Emmering et al. 1992) presented a torus outflow model. In their model, dusty molecular clouds are uplifted from the accretion disk and injected into outflows, and these moving optically thick clouds constitute a torus. Our scenario resembles theirs. But there are apparent differences. Firstly, the velocity of the outflows we studied is of order 10000 km/s, much faster than the torus outflows. Second, the clouds in BAL outflows are unlikely to be optically thick in dust opacity, since a large fraction of BAL quasars can be detected at UV band and the average amount of reddening in BAL quasars is only $E(B-V)\sim0.023$ (Reichard et al. 2003). Thirdly, the large fraction of BAL outflows is highly ionized. One solution to the discrepancies is that the torus structure is stratified. The dusty BAL outflow is perhaps the inner edge of the torus outflow of Elitzur \& Shlosman and is in an extreme ionization and dynamical state compared to the bulk torus. In this case, it is reasonable to speculate the existence of an outflow component with ionization and velocity in between. Recently, Zhang et al. (2013) found, in low redshift AGNs, an interesting connection between mid-infrared (MIR) emission and outflows in NLR, the velocity of which is of order several hundred kilometers per second (see also H\"{o}nig \etal\ 2013). The NLR outflows seem to fill the gap. However, a simple extrapolation from local AGNs may not be appropriate. Studies of the relationship between various outflow components (such as NLR, BAL, BEL) and MIR emission of $z\sim2$ quasars may help in uncovering the underlying physics. Before discussing the second mechanism, let us list several observational facts. The maximum velocity of a typical BAL outflow is about 10000 km/s. The typical width of BELs is about 4000 km/s and BAL outflows are co-spatial with or outside of BEL regions. It means that the typical Keplerian velocity at BAL outflow location is less than 4000 km/s. The last fact is the absence of evidence for outflows with similar velocity at galactic scale or in NLR. All of these facts together show that BAL outflows are strongly decelerated before reaching NLR and the gravity of their central black holes, however, is not responsible for the deceleration. Thus it is likely that outflows interact with surrounding medium. This medium is located between BLR and NLR and should be dense enough that it can effectively decelerate the outflows. Apparently, interaction with torus clouds is the most favorable. Recent hydrodynamical simulation (Wagner et al. 2013) showed that fast outflows can break dense clouds into diffuse warm filaments. This process makes more dust in the clouds exposed to the central UV source. As a consequence, the infrared emission increases and the outflows become dusty. Since a stronger outflow can ablate dense clouds more effectively, the dependence of NIR emission on outflow strength is yielded. In order to reach a high enough temperature to radiate in NIR band, the interaction has to occur at the innermost region of torus. In this scenario, outflows confine the geometry and subtending angle of dusty torus. It is in line with the suggestion that outflows blow away ambient dusty gas and transform a buried quasar into normal quasar phase (e.g. Sanders et al. 1988; Rupke \& Veilleux 2013). One problem of this scenario is that the interaction timescale (see e.g. Wagner et al. 2013) is much shorter than the quasar lifetime. The correlation might disappear after outflows blow away all of clouds in the outflow direction.
14
3
1403.3166
We have investigated a sample of 2099 broad absorption line (BAL) quasars with z = 1.7-2.2 built from the Sloan Digital Sky Survey Data Release Seven and the Wide-field Infrared Survey. This sample is collected from two BAL quasar samples in the literature and is refined by our new algorithm. Correlations of outflow velocity and strength with a hot dust indicator (β<SUB>NIR</SUB>) and other quasar physical parameters—such as an Eddington ratio, luminosity, and a UV continuum slope—are explored in order to figure out which parameters drive outflows. Here β<SUB>NIR</SUB> is the near-infrared continuum slope, which is a good indicator of the amount of hot dust emission relative to the accretion disk emission. We confirm previous findings that outflow properties moderately or weakly depend on the Eddington ratio, UV slope, and luminosity. For the first time, we report moderate and significant correlations of outflow strength and velocity with β<SUB>NIR</SUB> in BAL quasars. It is consistent with the behavior of blueshifted broad emission lines in non-BAL quasars. The statistical analysis and composite spectra study both reveal that outflow strength and velocity are more strongly correlated with β<SUB>NIR</SUB> than the Eddington ratio, luminosity, and UV slope. In particular, the composites show that the entire C IV absorption profile shifts blueward and broadens as β<SUB>NIR</SUB> increases, while the Eddington ratio and UV slope only affect the high and low velocity part of outflows, respectively. We discuss several potential processes and suggest that the dusty outflow scenario, i.e., that dust is intrinsic to outflows and may contribute to the outflow acceleration, is most likely.
false
[ "Infrared Survey", "outflow velocity", "outflow strength", "outflows", "UV slope", "the Sloan Digital Sky Survey Data Release Seven", "hot dust emission", "Seven", "non-BAL quasars", "BAL quasars", "blueshifted broad emission lines", "velocity", "SUB", "other quasar physical parameters", "Eddington", "NIR</SUB", "the Wide-field Infrared Survey", "luminosity", "non-BAL", "strength" ]
16.036875
7.314981
119
745633
[ "Prinja, R. K.", "Urbaneja, M. A." ]
2014MNRAS.440.2684P
[ "Rummaging inside the Eskimo's parka: variable asymmetric planetary nebula fast wind and a binary nucleus?" ]
7
[ "Department of Physics &amp; Astronomy, University College London, Gower Street, London WC1E 6BT, UK", "Institute for Astro- and Particle Physics, University of Innsbruck, Technikerstr. 25/8, A-6020 Inssbruck, Austria; Institute for Astronomy, University of Hawaii, 2680 Woodlawn Drive, Honolulu, HI 96822, USA" ]
[ "2015ApJ...800....8M", "2015RMxAA..51..221Z", "2016NewA...45....7G", "2017Ap.....60..497K", "2019ApJ...884..134G", "2019PASA...36...18M", "2021ApJ...909...44G" ]
[ "astronomy" ]
7
[ "stars: evolution", "stars: individual: NGC 2392", "stars: mass-loss", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1987AJ.....93..968R", "1991IAUS..145..375M", "1994ApJ...435..416B", "1995A&AS..110..353P", "1995ApJ...438..376M", "1995ApJ...439..875H", "1996ApJS..103..475F", "1997MNRAS.291..658D", "1998ApJ...496..407H", "2003Msngr.114...20M", "2003SPIE.4843..425M", "2004ApJ...602L..93D", "2006IAUS..234..153G", "2011MNRAS.417.2440H", "2012A&A...542A..45K", "2012AJ....144...58K", "2012ApJ...759L..28P", "2012ApJ...761..172G", "2012IAUS..282..470D", "2012IAUS..283..436M", "2012MNRAS.422.3142P", "2013A&A...553A.126G" ]
[ "10.1093/mnras/stu452", "10.48550/arXiv.1403.1480" ]
1403
1403.1480_arXiv.txt
Planetary nebulae (PNe) are a key pathway in the evolution of low to intermediate mass stars, and their central stars are the immediate precursors of white dwarfs. Studies of PN central stars (herein CSPNe) are motivated by: The desire to understand the origin of the rich variety of PN morphologies; to establish the mass-loss process via fast winds driven by radiation pressure by spectral lines; and to secure fundamental stellar parameters that can test post-AGB stellar evolution models. Time-series spectroscopy is an important diagnostic tool in developing our understanding of CSPNe. Recently, far-UV and UV datasets have revealed signatures of large-scale wind structures and evidence for modulated temporal behaviour that may provide a handle on the central star rotation rates (e.g. Prinja et al. 2012a, 2012b). Similarities between the wind properties of H-rich CSPNe and those of massive Population I OB stars (which also have line-driven winds) suggest that instabilities in variable fast winds may result in shock heated gas which emits X-rays in the central cavities of PNe (e.g. Guerrero 2006; Kastner et al. 2012). In the optical waveband, time-series data are requisite for establishing systematic radial velocity shifts in CSPNe absorption lines. De Marco et al. (2004) have for example conducted a radial velocity survey of 11 CSPNe to provide constraints on the binary properties of the parent AGB population and thus the extent to which binarity may play a causal role in shaping non-spherical nebulae. In this paper we present time-series optical spectra of the central star of NGC~2392 (Eskimo nebula). Our study is motivated by several interesting characteristics, discrepancies and scenarios for this PN: (i) The central star of NGC~2392 exhibits high He, N and low C, O abundances suggesting that the photosphere has been processed (M{\'e}ndez et al. 2012). A possible scenario is that the abnormal central star abundances are due to a common-envelope evolutionary phase thus implying a close binary companion; (ii) Danehkar et al. (2011) employ photoionization models of high excitation PN emission lines to argue that NGC~2392 has a hot white dwarf ($\sim$ 1 M$_\odot$) companion; (iii) Detailed kinematic modelling of the (Eskimo) nebula by Garcia-Diaz et al. (2012) supports a near-pole orientation, complex nebula morphology with multiple kinematic components, and an evolution path that may invoke a common-envelope binary; (iv) The extended and point X-ray emission from NGC~2392 (e.g. Kastner et al. 2012) is not entirely consistent with the predicted thermal energy converted from the kinetic energy of the fast wind. Additional coronal energy from a binary companion may explain the observed high X-ray temperatures. Despite all the implications of the above studies, there is no definitive evidence so far of a binary nucleus in NGC~2392, and the time-variable and geometric characteristics of its fast wind are not established. In this study we present the analysis of high-resolution optical time-series datasets secured over two epochs in 2006 and 2010 using the 3.6m ESO and Canada-France-Hawaii (CFHT) observatories. Our goal is to investigate for the first time $\sim$ hourly changes in the fast wind of NGC~2392 and fluctuations close to the surface of the central star. We characterise here evidence for evolving structure in the outflow and indications of radial velocity changes in deep-seated absorption lines.
We have provided evidence and demonstrated that in two independently secured optical spectroscopic time-series separated by $\sim$ 3 years, both datasets reveal for the central star of NGC~2392; (i) stochastic variations in the fast wind-formed recombination lines on timescales down to $\sim$ 30 min., (ii) changes in the overall morphology of the He{\sc i} and He{\sc ii} line profiles that cannot be accounted for by 1-D line-synthesis predictions for a spherically homogeneous wind, (iii) radial velocity shifts of semi-amplitude $\sim$ 10 km s$^{-1}$ in N{\sc iv} $\lambda$6380.8 (and tentatively in H$\epsilon$) that show maximum Fourier power spectra signal at $\sim$ 0.123-d in ESO (2006) and CFHT (2010) data. Our detection of radial velocity motion of the central star in NGC~2392 is obviously tentative. Cross-correlation with the majority of stellar absorption lines in the optical range is not fruitful since most of features are `contaminated' by the imprints of stellar wind variability which we have established here. In advance of more definitive optical spectroscopy and photometry being secured, we can only speculate on the potential binary components in NGC~2392: For an assumed circular orbit and line-of-sight inclination $\sim$ $10\degr$ (e.g. Garcia-Diaz et al. 2012); semi-amplitude of 10 km s$^{-1}$ and period = 0.123-d (Figs. 8 and 9); assumed (primary) central star mass = 0.6 M$_\odot$, the implied dynamical mass of the secondary in circular Keplerian orbit is $\sim$ 0.1 M$_\odot$. Such a low-mass companion would for example correspond to a late M-dwarf of $T_{\rm eff}$ $\sim$ 2000 - 3000K. \subsection{Constraints from UV line profile morphologies} The morphologies of the FUV and UV lines are complex in NGC~2392 and provide additional signatures for an asymmetric geometry. The CMFGEN line-synthesis models (Sect. 3) do not provide consistent matches to Doppler widths, and absorption and emission strengths across all the UV ion stages observed in the fast wind. A selection of UV wind line profiles in NGC~2392 is presented in Fig. 9, ranging from O{\sc vi} $\lambda$1031.9 and S{\sc vi} $\lambda$944.5, to P{\sc v} $\lambda$1118.0 and N{\sc v} $\lambda$1238.8, and C{\sc iv} $\lambda$1548.2 and Mg{\sc ii} $\lambda$2795.5. (The data have been retrieved from the {\it FUSE} and {\it IUE} archives.) There are some key points to note in Fig 9: (i) All the wind lines get weaker with increasing outflow velocity. It may be that the wind plasma is shifting to a very high ionization state (beyond O{\sc vi}) as it travels to larger radii. Alternatively, the line shapes in Fig. 9 could be an indication that the fast wind of NGC~2392 is moving out of the line-of-sight, somewhat as may be expected for a polar, high-latitude wind in an asymmetric geometry; (ii) The low excitation Mg{\sc ii} line is very narrow ($\sim$ 100 km s$^{-1}$; as is Si{\sc iii} $\lambda$1206.5) and consistent with a low-velocity equatorial wind (see e.g. Bjorkman et al. 1994; Massa 1995); (iii) The presence of P{\sc v} most likely indicates that the wind is optically very thick, since phosphorous has a low cosmic abundance and this line would otherwise not be so clearly detected. However C{\sc iv} and N{\sc v} are weak at intermediate velocities ($\sim$ 200 to 300 km s$^{-1}$). An optically very thick wind that causes weak absorption can arise in a scenario where the wind is not covering the entire stellar disk, as may be expected from a polar wind. We conclude that UV lines provide evidence for an asymmetric, two-component outflow in NGC~2392, where high-speed high-ionization gas forms preferentially in the polar region. Slower, low ionization material is then confined primarily to a cooler equatorial component of the outflow. \begin{figure} \includegraphics[scale=0.25]{ngc2392_prinja_urbaneja_fig9.pdf} \caption{The fast wind in NGC~2392 as revealed by the FUV and UV resonance line profiles of O{\sc vi} $\lambda$1031.9, S{\sc vi} $\lambda$944.5, P{\sc v} $\lambda$1118.0, N{\sc v} $\lambda$1238.8, C{\sc iv} $\lambda$1548.2, Mg{\sc ii} $\lambda$2795.5. } \end{figure}
14
3
1403.1480
We report on high-resolution optical time series spectroscopy of the central star of the `Eskimo' planetary nebula NGC 2392. Data sets were secured with the European Southern Observatory (ESO) 2.3 m in 2006 March and Canada-France-Hawaii Telescope (CFHT) 3.6 m in 2010 March to diagnose the fast wind and photospheric properties of the central star. The He I and He II recombination lines reveal evidence for clumping and temporal structures in the fast wind that are erratically variable on time-scales down to ∼30 min (i.e. comparable to the characteristic wind flow time). We highlight changes in the overall morphology of the wind lines that cannot plausibly be explained by line-synthesis model predictions with a spherically homogeneous wind. Additionally, we present evidence that the ultravoilet line profile morphologies support the notion of a high-speed, high-ionization polar wind in NGC 2392. Analyses of deep-seated, near-photospheric absorption lines reveals evidence for low-amplitude radial velocity shifts. Fourier analysis points tentatively to an ∼0.12-d modulation in the radial velocities, independently evident in the ESO and CFHT data. We conclude that the overall spectroscopic properties support the notion of a (high-inclination) binary nucleus in NGC 2392 and an asymmetric fast wind.
false
[ "NGC", "the characteristic wind flow time", "an asymmetric fast wind", "∼30 min", "the fast wind", "photospheric properties", "the wind lines", "high-resolution optical time series spectroscopy", "Eskimo", "Data sets", "time-scales", "ESO", "the central star", "line-synthesis model predictions", "a spherically homogeneous wind", "binary nucleus", "the ultravoilet line profile morphologies", "NGC 2392", "a high-speed, high-ionization polar wind", "evidence" ]
10.07984
10.293113
151
746752
[ "Bartos, I.", "Veres, P.", "Nieto, D.", "Connaughton, V.", "Humensky, B.", "Hurley, K.", "Márka, S.", "Mészáros, P.", "Mukherjee, R.", "O'Brien, P.", "Osborne, J. P." ]
2014MNRAS.443..738B
[ "Cherenkov Telescope Array is well suited to follow up gravitational-wave transients" ]
25
[ "Department of Physics, Columbia University, New York, NY 10027, USA; Columbia Astrophysics Laboratory, Columbia University, New York, NY 10027, USA", "Department of Astronomy and Astrophysics, Pennsylvania State University, University Park, PA 16802, USA; George Washington University, Department of Physics, 725 21st St, NW, Washington, DC 20052, USA", "Columbia Astrophysics Laboratory, Columbia University, New York, NY 10027, USA", "Center for Space Plasma and Aeronomic Research (CSPAR), University of Alabama in Huntsville, Huntsville, AL 35899, USA", "Columbia Astrophysics Laboratory, Columbia University, New York, NY 10027, USA", "Space Sciences Laboratory, University of California-Berkeley, Berkeley, CA 94720, USA", "Department of Physics, Columbia University, New York, NY 10027, USA; Columbia Astrophysics Laboratory, Columbia University, New York, NY 10027, USA", "Department of Astronomy and Astrophysics, Pennsylvania State University, University Park, PA 16802, USA", "Department of Physics and Astronomy, Barnard College, Columbia University, New York, NY 10027, USA", "Space Research Centre, Department of Physics and Astronomy, University of Leicester, Leicester LE1 7RH, UK", "Space Research Centre, Department of Physics and Astronomy, University of Leicester, Leicester LE1 7RH, UK" ]
[ "2014ApJ...787..168V", "2015ApJ...805...82M", "2015ICRC...34..763B", "2015MNRAS.450.1430H", "2015PASA...32...46H", "2016ApJ...816...61B", "2016ApJ...828L...4Z", "2016MNRAS.459..121C", "2016sf2a.conf...31S", "2017JPhCS.888a2001B", "2018JCAP...05..056P", "2018MNRAS.477..639B", "2019ICRC...36..790S", "2019MNRAS.490.3476B", "2019NPPP..306...20H", "2019NPPP..306...69D", "2019PASA...36...19D", "2019arXiv190105486C", "2019scta.book.....C", "2020PhRvD.101l3008G", "2021ApJ...909..207R", "2021ApJ...918...66A", "2021PAN....84.1037P", "2022icrc.confE.998P", "2023MNRAS.522.5690M" ]
[ "astronomy" ]
7
[ "gravitational waves", "methods: observational", "gamma-ray burst: general", "Astrophysics - High Energy Astrophysical Phenomena" ]
[ "1976PhFl...19.1130B", "1984SvAL...10..177B", "1986ApJ...308L..43P", "1989Natur.340..126E", "1996ApJ...457..253D", "1997PhRvL..78.2292W", "1998ApJ...496L...1R", "1998ApJ...507L..59L", "1999ApJ...520..641S", "2001ApJ...548..787S", "2001ApJ...552L..35Z", "2001ApJ...559..110Z", "2002ApJ...568..820G", "2003ASPC..289...21E", "2004AdSpR..34.2657B", "2005SSRv..120..143B", "2005astro.ph.10256K", "2006A&A...453..823G", "2006ApJ...636..777A", "2006ApJ...641L...9A", "2006ApJ...643..266N", "2006ApJ...648..774S", "2006ApJ...650..281N", "2006PhRvD..74h2004C", "2007ApJ...668..392G", "2007NJPh....9...17L", "2007PhR...442..166N", "2007PhRvD..75f2004R", "2008ApJ...688L..61V", "2008CQGra..25r4034K", "2008MNRAS.390..781M", "2009ApJ...697.1071A", "2009ApJ...702..791M", "2009ApJ...703..675N", "2009ApJ...706L.138A", "2009CQGra..26k4007R", "2009CQGra..26t4010V", "2009IEEEP..97.1431D", "2009NJPh...11l3006F", "2009Natur.462..331A", "2009arXiv0902.1527B", "2010ApJ...709L.146D", "2010ApJ...712..558A", "2010ApJ...716.1178A", "2010ApJ...723.1711T", "2010ApJS..188..405A", "2010CQGra..27q3001A", "2010MNRAS.403..926G", "2010MNRAS.406.2650M", "2010MNRAS.407.1033B", "2010MNRAS.409..226K", "2010PhRvD..81h2001W", "2011ApJ...729..114A", "2011ApJ...730..141Z", "2011ApJ...733L..37V", "2011ApJ...739...99N", "2011ApJ...743...62A", "2011CQGra..28h5016W", "2011CQGra..28j5021F", "2011CQGra..28l5023S", "2011ExA....32..193A", "2011IJMPD..20.1755K", "2011MNRAS.410.2556D", "2011MNRAS.415.1663T", "2011Natur.478...82N", "2011PhRvD..83j2001K", "2011PhRvD..84j4020V", "2011PhRvL.107y1101B", "2011arXiv1109.5680B", "2012A&A...539A.124L", "2012A&A...541A.155A", "2012ApJ...746...48M", "2012ApJ...748..136C", "2012ApJ...752...29H", "2012ApJ...753L..31A", "2012ApJ...759...22K", "2012ApJS..203...28E", "2012MNRAS.421..525H", "2012MNRAS.425.2668C", "2012NIMPA.692...72D", "2012PhRvD..85h2002A", "2012PhRvD..85j4045V", "2012PhRvL.108w1101H", "2012RAA....12.1139M", "2013APh....43....3A", "2013APh....43...19H", "2013APh....43..134M", "2013APh....43..171B", "2013APh....43..189D", "2013APh....43..252I", "2013APh....43..317D", "2013ApJ...764...94V", "2013ApJ...764..179B", "2013ApJ...767..124N", "2013ApJ...773L..20L", "2013ApJ...774...25K", "2013ApJ...776...18F", "2013ApJS..209...11A", "2013CQGra..30l3001B", "2013MNRAS.430.2121P", "2013MNRAS.435..502F", "2013Natur.500..547T", "2013PhRvL.111m1102M", "2013RvMP...85.1401A", "2014ARA&A..52...43B", "2014ApJ...787..168V", "2014ApJ...789L...5K", "2014ApJ...795..105S", "2014ApJS..211....7A", "2014MNRAS.437..649S", "2014Sci...343...42A", "2018LRR....21....3A" ]
[ "10.1093/mnras/stu1205", "10.48550/arXiv.1403.6119" ]
1403
1403.6119_arXiv.txt
Observing the electromagnetic counterparts of the first detected gravitational-wave (GW) signals is one of the major goals of astronomy for the near future \citep{2009arXiv0902.1527B,2012ApJ...759...22K}. Electromagnetic counterparts would greatly increase our confidence in the first detection, and could revolutionize our understanding of some cosmic phenomena (e.g., \citealt{2012A&A...541A.155A,2012ApJS..203...28E,2013CQGra..30l3001B,2013arXiv1310.2314T}). One of the most anticipated cosmic phenomena that is expected to result in the first GW detections is the merger of two compact stellar-mass objects, which are either neutron stars or black holes \citep{2012PhRvD..85h2002A}. Of special interest are binaries that consist of at least one neutron star, which can also produce electromagnetic radiation \citep{1984SvAL...10..177B,1986ApJ...308L..43P,1989Natur.340..126E,2007NJPh....9...17L,2013CQGra..30l3001B} as well as other messengers, such as cosmic rays or neutrinos \citep{1997PhRvL..78.2292W,PhysRevLett.108.231101,2012ApJ...752...29H,PhysRevLett.107.251101,2012arXiv1203.5192A,2013PhRvL.111m1102M}. Several promising emission processes have been suggested that would accompany compact binary mergers. First of all, short gamma-ray bursts (GRBs; \citealt{2013APh....43..134M}) are thought to originate from these mergers. Gamma rays can be produced in the outflows driven by an accreting black hole that forms in the merger (e.g., \citealt{2007PhR...442..166N}). The afterglows of some of these short GRBs, produced by the interaction of the outflow with the ambient medium, represent an additional electromagnetic counterpart \citep{1999ApJ...520..641S,2011ApJ...733L..37V}. Further, energetic, sub or mildly relativistic outflows launched by a binary merger can also interact with the surrounding medium, producing quasi-isotropic emission in the radio band over a period of more than a year \citep{2011Natur.478...82N,2013MNRAS.430.2121P}. The same outflow can also undergo r-process nucleosynthesis during its expansion, resulting in near-infrared--infrared radiation, called a kilonova (also called macronova; \citealt{1998ApJ...507L..59L,2005astro.ph.10256K,2008MNRAS.390..781M,2010MNRAS.406.2650M,2013ApJ...774...25K,2013MNRAS.435..502F,tanvir2013kilonova,2013arXiv1311.2603B}). With the available observational capabilities, a large fraction of the produced electromagnetic counterparts may only be detectable, if detectable at all, with follow-up observations guided by GW triggers \citep{2008CQGra..25r4034K,2011ApJ...733L..37V,2012ApJ...746...48M,2012A&A...539A.124L,2012ApJS..203...28E,2013CQGra..30l3001B}. The direction reconstruction of GW detectors, however, is limited to $\gg1$\,deg$^2$ \citep{2013arXiv1304.0670L}. This substantially reduces the feasibility of many electromagnetic follow-up efforts, given the limited field of view of the most sensitive telescopes, and the limited sensitivity of larger-field-of-view telescopes. Nevertheless, a number of telescopes may be competitive at following up GW triggers with very low latency with strategies optimized to cover a significant fraction of the GW sky area. These include moderate-aperture telescopes with large field of view such as the BlackGEM Array\footnote{https://www.astro.ru.nl/wiki/research/blackgemarray} and the Ground Wide Angle Cameras (GWAC; \citealt{2009AIPC.1133...25G}) that will be dedicated follow-up operations. Highly sensitive instruments with limited field of view, such as Swift, may also be promising follow-up facilities \citep{2012ApJS..203...28E}. Another interesting direction is the Ultra Fast Flash Observatory \citep{2012SPIE.8443E..0IP}, which will have a large field-of-view x-ray detector as well as sub-second optical follow-up capability. Following up the first GW observations, probably around 2016-2018 \citep{2013arXiv1304.0670L}, will be particularly challenging. At this time, given the staged schedule of construction and commissioning of GW detectors, it is possible that direction reconstruction of the first detections will mainly rely on a two-detector network (or a three-detector network with highly non-uniform sensitivity; \citealt{2013arXiv1304.0670L}). This will substantially decrease the location accuracy of GW measurements, necessitating electromagnetic follow-ups with large, $100-1000$\,deg$^2$, search areas at high sensitivity. Further, the direction of a GW event will typically be localized to multiple, disjoint sky regions at potentially distant parts of the sky, requiring follow-up observations to cover these separate sky regions. Radio follow-up observations, e.g., with the Square Kilometre Array (SKA; \citealt{2003ASPC..289...21E}) or LOFAR \citep{2009IEEEP..97.1431D}, are another interesting alternative, given the expected long-lived radio emission following the binary merger \cite{2013MNRAS.430.2121P}. In this paper we propose and investigate the possibility for large-field-of-view electromagnetic follow-up observations of GW event candidates, using the Cherenkov Telescope Array (CTA; \citealt{2011ExA....32..193A}). CTA is well suited for GW-follow-up observations for multiple reasons: \begin{enumerate} \item \emph{Field-of-view:} CTA will be capable of monitoring a large sky area via survey mode operation (either by pointing its telescopes in different directions, or by rapidly scanning a set of consecutive directions). It will be able to monitor the $\sim$$1000$\,deg$^2$ area necessary for early GW triggers for which only an incomplete GW detector network is available. This survey mode will also be useful for later GW observations: since localization becomes less efficient with, e.g., decreasing signal-to-noise ratio (SNR), a significant fraction of GW event candidates will have large error regions even when more than two GW detectors are available. \item \emph{Coincident observational schedule:} CTA is expected to begin partial operation around 2017, therefore it will probably be available to follow up the first GW detections. The anticipated completion of CTA is around 2020. \item \emph{Rapid response:} CTA has the capability to respond to target-of-opportunity requests and start monitoring the selected sky area within $\sim$\,$30$\,sec \citep{2013APh....43..317D}. This is important given the limited duration ($\lesssim1000$\,s; Section \ref{section:CTAsensitivity}) of high-energy photon emission connected with GRBs. The sensitivity of CTA to GRBs will be mainly determined by its so-called Large Size Telescopes (LST; \citealt{2013APh....43....3A}), which are capable of the fastest response ($180^\circ$ slewing in less than 20\,s; \citealt{Inoue+13GRBCTA}). \end{enumerate} In the following we discuss these points further in detail. The use of CTA to follow up GW event candidates has been previously suggested by \cite{2013APh....43..189D}. In this paper, we explore in detail the follow-up of GW event candidates by CTA, and the particular advantage of CTA in following-up poorly localized signals. The paper is organized as follows. In Section \ref{section:GW} we discuss the expected sensitivity and localization capability of advanced GW detectors. In Section \ref{section:CTAsensitivity} we estimate the sensitivity of CTA for detecting short GRBs with known directions, exploring multiple emission models focusing on the distances relevant for GW detection. In Section \ref{section:CTAsurvey} we describe the sensitivity of CTA in survey mode, focusing on directional uncertainties relevant for GW searches. Section \ref{section:GBM} discusses the role of satellite-based GRB detectors in adding information to the follow-up process. Finally, Section \ref{section:conclusion} summarizes our results and presents our conclusions.
\label{section:conclusion} We explored the feasibility of following up GW events with CTA. We focused on the scenario in which the GW event is poorly localized, necessitating follow-up observations to cover up to $\sim$$1000$\,deg$^2$ of sky area. Limited localization can emerge from various detection scenarios. In the early advanced GW detector era, one can expect only the two LIGO observatories to operate at high sensitivity, and direction reconstruction with two detectors is limited. But even with further GW detectors in operation, GW event candidates with relatively low signal-to-noise ratios will also have poorly constrained directions of origin, therefore requiring follow-up over a larger sky area. We based our study on short GRBs, assuming that they originate from compact binary mergers, which are considered the most promising sources for the first GW detections. While various follow-up observations (e.g., optical/infrared) will be difficult to carry out over larger ($\gg 100$\,deg$^2$) sky areas with the desired sensitivity, we find that CTA may be capable of efficiently detecting late-time high-energy gamma-ray emission from short GRBs. To estimate their detectability, we extrapolated the energy spectrum observed by Fermi-LAT to $\gtrsim 50$\,GeV where CTA becomes sensitive. Currently it is unclear, due to the lack of observations, whether short-GRB spectra extend into this range, and to how high an energy. We considered different cutoff energies (from 50\,GeV to 1\,TeV), as well as multiple GRB emission scenarios, to investigate the sensitivity of CTA for these different cases. Our results show that short GRBs with high-energy emission extending up to $\sim100$\,GeV can be detectable via CTA, even if CTA needs to survey a sky area of $\sim$$1000$\,deg$^2$ and if CTA observations are delayed by $\sim 100$\,s following the onset of gamma-ray emission. Detection with lower energy cutoffs is also promising, although may require a dense circumburst medium, faster GW event reconstruction, smaller sky area, or closer source. For comparison we considered a $\sim$$200$\,deg$^2$ sky area that can be achieved for some events with stronger GW emission, or if we restrict our search to a fraction of the sky area with the highest-probability directions. For a $\sim$$200$\,deg$^2$ sky area, we find that GRBs even with cutoffs somewhat below $100$\,GeV can be detectable, although for a $\sim50$\,GeV cutoff one requires faster response than $\sim 100$\,s. Many of the events detected by both GW facilities and CTA will also likely be observed by GRB satellites. For observations with low latency, as in the case of Fermi-GBM and Swift-BAT, the localization of the GRB can significantly reduce the sky area CTA needs to cover in order to find the source. The detection of MeV emission can also be important in mapping the connection between GW and electromagnetic emission within a broad energy range. We estimated the rate of events that can be jointly detected by CTA and GW observatories, consider a characteristic short-GRB rate of $10$\,Gpc$^{-3}$yr$^{-1}$, and a fiducial GW horizon distance of of $300$\,Mpc. With $\sim$$11\%$ duty cycle for CTA we find a limited detection rate of $\sim$$0.03$\,yr$^{-1}$. A further decrease of $\sim 50\%$ is expected from CTA being able to observe only for source elevations $\gtrsim 30^\circ$ above the horizon. This number, nevertheless, can increase if sub-threshold GW events with lower SNRs are followed up, or if a sub-population of short GRBs originate from black hole-neutron star mergers, which can be detected by GW observatories from larger distances. The estimate nevertheless indicates that a joint detection may require an extended period of operation. Overall, we find that CTA is well suited to perform follow-up observations of GW events, even those with limited source localization. It can, therefore, be important to carry out a more detailed investigation of the possible follow-up observation strategies with CTA, and the expected joint sensitivity, beginning as early as in the installation phase. It will also be important to further our understanding of the phenomenology of GRB emission at $\gg$\,GeV energies.
14
3
1403.6119
The first gravitational-wave (GW) observations will greatly benefit from the detection of coincident electromagnetic counterparts. Electromagnetic follow-ups will nevertheless be challenging for GWs with poorly reconstructed directions. GW source localization can be inefficient (i) if only two GW observatories are in operation; (ii) if the detectors' sensitivities are highly non-uniform; (iii) for events near the detectors' horizon distance. For these events, follow-up observations will need to cover 100-1000 deg<SUP>-2</SUP> of the sky over a limited period of time, reducing the list of suitable telescopes. We demonstrate that the Cherenkov Telescope Array (CTA) will be capable of following up GW event candidates over the required large sky area with sufficient sensitivity to detect short gamma-ray bursts, which are thought to originate from compact binary mergers, out to the horizon distance of advanced LIGO/Virgo. CTA can therefore be invaluable starting with the first multimessenger detections, even with poorly reconstructed GW source directions. This scenario also provides a further scientific incentive for GW observatories to further decrease the delay of their event reconstruction.
false
[ "GW observatories", "suitable telescopes", "compact binary mergers", "GW source localization", "sufficient sensitivity", "GW", "events", "LIGO", "advanced LIGO/Virgo", "coincident electromagnetic counterparts", "poorly reconstructed GW source directions", "short gamma-ray bursts", "the horizon distance", "the required large sky area", "time", "the detectors horizon distance", "their event reconstruction", "i", "operation", "poorly reconstructed directions" ]
7.216968
1.671919
-1
746197
[ "McKernan, B.", "Ford, K. E. S.", "Kocsis, B.", "Lyra, W.", "Winter, L. M." ]
2014MNRAS.441..900M
[ "Intermediate-mass black holes in AGN discs - II. Model predictions and observational constraints" ]
131
[ "Department of Science, Borough of Manhattan Community College, City University of New York, New York, NY 10007, USA; Department of Astrophysics, American Museum of Natural History, New York, NY 10024, USA; Graduate Center, City University of New York, 365 5th Avenue, New York, NY 10016, USA; Kavli Institute for Theoretical Physics, UC Santa Barbara, CA 93106, USA", "Department of Science, Borough of Manhattan Community College, City University of New York, New York, NY 10007, USA; Department of Astrophysics, American Museum of Natural History, New York, NY 10024, USA; Graduate Center, City University of New York, 365 5th Avenue, New York, NY 10016, USA; Kavli Institute for Theoretical Physics, UC Santa Barbara, CA 93106, USA", "Harvard-Smithsonian Center for Astrophysics, 60 Garden St, Cambridge, MA 02138, USA; Institute for Advanced Study, Einstein Drive, Princeton, NJ 08540, USA", "Jet Propulsion Laboratory, California Institute of Technology, 4800 Oak Grove Dr., Pasadena, CA 91109, USA; California Institute of Technology, Division of Geological &amp; Planetary Sciences, 1200 E. California Blvd, Pasadena, CA 91125, USA", "Center for Astrophysics &amp; Space Astronomy, University of Colorado, Boulder, CO 80303, USA" ]
[ "2014MNRAS.445L..74M", "2015A&A...576A..29I", "2015MNRAS.451.1341L", "2015MNRAS.452L...1M", "2016ApJ...819L..17B", "2017IJMPD..2630021M", "2017MNRAS.464..946S", "2017MNRAS.470.4112G", "2017NatCo...8..831B", "2017mbhe.confE..51K", "2018ApJ...856...47T", "2018ApJ...856...92F", "2018ApJ...856..113T", "2018ApJ...860..129S", "2018ApJ...864...27S", "2018ApJ...866...66M", "2018ApJ...867..119F", "2018FrASS...5...19B", "2018MNRAS.474.5672L", "2018MNRAS.480.4468R", "2018MNRAS.480.5160F", "2018PhRvD..97l3003K", "2018arXiv180909130M", "2019ApJ...878...85S", "2019ApJ...883..139S", "2019ApJ...884L..50M", "2019ApJ...886...25B", "2019Atoms...7...26B", "2019BAAS...51c.123B", "2019BAAS...51c.175B", "2019BAAS...51c.247F", "2019IAUS..346..397M", "2019MNRAS.486.2754D", "2019MNRAS.488...47F", "2019MNRAS.490L..42F", "2019PhRvL.123r1101Y", "2020ApJ...892..130D", "2020ApJ...894..133A", "2020ApJ...896..138Y", "2020ApJ...897...46F", "2020ApJ...897..100V", "2020ApJ...897..142S", "2020ApJ...898...25T", "2020ApJ...899...26T", "2020ApJ...899..149R", "2020ApJ...900L..13A", "2020ApJ...903..133S", "2020ApJ...904L..13R", "2020CQGra..37u5011A", "2020MNRAS.493.3732D", "2020MNRAS.494.1203M", "2020MNRAS.495..536E", "2020MNRAS.498.2339R", "2020MNRAS.498.3452D", "2020MNRAS.498.4088M", "2020MNRAS.499.2608F", "2020PhRvD.101f3002T", "2020PhRvD.102d3015A", "2020PhRvL.124y1102G", "2021A&A...650A.107M", "2021ApJ...906L..11Z", "2021ApJ...907L..48V", "2021ApJ...908L..27S", "2021ApJ...910...94C", "2021ApJ...910..152Z", "2021ApJ...911..124L", "2021ApJ...914L..34P", "2021ApJ...919..140S", "2021ApJ...921L..43Z", "2021ApJ...923..173P", "2021ExA....51.1427A", "2021ExA....51.1441B", "2021MNRAS.500.4307F", "2021MNRAS.501.3540D", "2021NatAs...5..749G", "2021PhLB..82236654F", "2021PhRvD.103j3018P", "2021PhRvD.104h2003O", "2021arXiv211212481C", "2022A&A...660A..26B", "2022ApJ...928..191G", "2022ApJ...929..133J", "2022ApJ...931...82V", "2022ApJ...933..170F", "2022ApJ...934..154L", "2022ApJ...935L..20Z", "2022ApJ...939L..23C", "2022ApJ...940..155D", "2022ApJ...940L..44R", "2022ExA....54...23P", "2022Galax..10...76S", "2022JCAP...09..034E", "2022MNRAS.514.3886M", "2022MNRAS.514.4102M", "2022MNRAS.515.3299G", "2022MNRAS.515.5106W", "2022MNRAS.516.5309V", "2022MNRAS.517.1602L", "2022MNRAS.517.5827F", "2022arXiv220406002M", "2023ARA&A..61..517L", "2023ApJ...942...99G", "2023ApJ...944...44K", "2023ApJ...944L..42L", "2023ApJ...948..120C", "2023ApJ...954...84W", "2023LRR....26....2A", "2023MNRAS.520.4502W", "2023MNRAS.521.2577P", "2023MNRAS.522..319C", "2023MNRAS.522.5393N", "2023MNRAS.524..940Z", "2023MNRAS.525..561E", "2023MNRAS.525.4202G", "2023Univ....9..138A", "2023arXiv231105393L", "2023arXiv231112118A", "2023arXiv231115778C", "2024A&A...685A..51V", "2024ApJ...962..143Q", "2024ApJ...967...88C", "2024MNRAS.527.6076R", "2024MNRAS.52710448R", "2024MNRAS.528.4958W", "2024MNRAS.528L.127W", "2024MNRAS.531.1409P", "2024MNRAS.531.3479M", "2024MNRAS.tmp.1420W", "2024PhRvD.109b4024M", "2024arXiv240408138D", "2024arXiv240502855L" ]
[ "astronomy" ]
14
[ "accretion accretion discs", "planets and satellites: formation", "planet-disc interactions", "protoplanetary discs", "galaxies: active", "galaxies: Seyfert", "Astrophysics - Astrophysics of Galaxies", "Astrophysics - High Energy Astrophysical Phenomena" ]
[ "1963PhRv..131..435P", "1964PhRv..136.1224P", "1973A&A....24..337S", "1983Ap&SS..95...11Z", "1986ApJ...309..846L", "1989MNRAS.238..729F", "1992A&A...262...97V", "1995ARA&A..33..581K", "1995ApJ...444..632Z", "1995MNRAS.277..758S", "1996ApJ...467L..77A", "1996Icar..124...62P", "1997MNRAS.286..513R", "1998ApJ...501L..29M", "1999ApJ...514..344B", "1999MNRAS.307...79I", "2001ApJ...555..775C", "2002ApJ...567L...9A", "2002MNRAS.330..232C", "2003AJ....126.1720S", "2003ApJ...588..494M", "2003ApJ...593.1116J", "2003MNRAS.340..411L", "2003MNRAS.341..501S", "2004ApJ...608..157M", "2004MNRAS.349L...7G", "2004PhRvD..69h2005B", "2004PhRvD..70l2002B", "2005ApJ...622L..93M", "2005ApJ...625..699M", "2005ApJS..160..401P", "2006A&A...455..773V", "2006ARA&A..44...49R", "2006ApJ...637...27K", "2006ApJ...648..411K", "2006Icar..181..587C", "2006MNRAS.365.1067C", "2007ApJ...670L.135E", "2007MNRAS.377.1647N", "2007MNRAS.379.1359M", "2007MNRAS.381.1280E", "2007PhRvD..76b2003K", "2008A&A...483..161T", "2008ARA&A..46..475H", "2008ASPC..399...20H", "2008ApJ...677.1184L", "2008ApJ...682.1134H", "2008ApJ...684..870K", "2008MNRAS.386..164P", "2008MNRAS.387.1063P", "2009MNRAS.395.2127O", "2010A&A...517A..47M", "2010A&A...524A..26C", "2010ApJ...711..796E", "2010ApJ...725..496N", "2010ApJS..187..416L", "2010MNRAS.407.2007C", "2010MNRAS.407.2399M", "2011ApJ...732...42A", "2011ApJ...734..111D", "2011ApJ...739...99N", "2011ApJ...742L..29R", "2011MNRAS.412.1591N", "2011MNRAS.413L..24M", "2011MNRAS.417L.103M", "2011PhRvD..83d4021H", "2011PhRvD..83d4030Y", "2011PhRvD..84b2002L", "2011PhRvD..84b4032K", "2011PhRvL.107q1103Y", "2012A&A...545A.127R", "2012ApJ...745..107W", "2012ApJ...750L..27S", "2012ApJ...751..104B", "2012ApJ...757...27A", "2012ApJ...759..118B", "2012ApJ...761...90G", "2012CQGra..29l4016A", "2012MNRAS.423.2633S", "2012MNRAS.423.3299N", "2012MNRAS.423L..65B", "2012MNRAS.425..460M", "2012MNRAS.427.2660K", "2012MNRAS.427.2680K", "2012PhRvD..85j4045V", "2012PhRvD..85l3005K", "2012PhRvD..86j4027M", "2013ApJ...762...35B", "2013ApJ...768..143Z", "2013ApJ...773..187N", "2013ApJ...774..144R", "2013ApJ...777...44J", "2013MNRAS.432.1468M", "2013MNRAS.436.2997D", "2013PhRvD..87d3004E", "2014ApJ...783...73F", "2014ApJ...783..134F", "2014ApJ...784...71S" ]
[ "10.1093/mnras/stu553", "10.48550/arXiv.1403.6433" ]
1403
1403.6433_arXiv.txt
\label{sec:intro} There is overwhelming observational evidence for supermassive black holes (SMBH;$>10^{6}M_{\odot}$) in the centers of galaxies \citep{b4} and stellar mass black holes ($<20M_{\odot}$) in our own Galaxy \citep{b2}. However, there is only fragmentary evidence for intermediate mass black holes (IMBHs, see \citealt{davis11} for the best candidate to date). IMBHs are thus a key missing component of our Universe. The standard model of IMBH production is in clusters \citep{b81}, however there are currently no undisputed cases of IMBHs in globular clusters \citep{b11}; IMBHs are hard to find. In \citet{b93} (Paper I), we described a model for the production and growth of IMBH seeds in disks around supermassive black holes. Our model grows IMBH seeds at super-Eddington rates in disks in active galactic nuclei (AGN). IMBH growth occurs both via collision of stars and compact objects at low relative velocity in disks (core accretion) and via gas accretion as the IMBH migrates within the disk. Our model is analagous to the growth of giant planets in protoplanetary disks \citep[e.g.][]{b89,b20} and grows IMBHs more efficiently than the standard model \citep{b81}. In this paper we outline the wide range of predicted observables that can reveal IMBHs in galactic nuclei throughout the local Universe. In Section~\S\ref{sec:observations} we discuss conditions in the AGN disk under which IMBHs can open a gap. In Section~\S\ref{sec:gap} we discuss observational consequences of gaps in the outer AGN disk. We draw a parallel between observational signatures of gapped protoplanetary disks and gaps carved out by IMBHs in AGN disks. In section~\S\ref{sec:feka} we outline the effect of a gap-opening IMBH in the inner AGN disk on the broad component of the FeK$\alpha$ line, which yields signatures that allow us to follow the final stage of mergers and provides advance warning of gravitational wave outbursts. In Section~\S\ref{sec:disk_nogap} we discuss the signatures of accreting IMBHs in AGN disks. In Section~\S\ref{sec:indep} we discuss occasional signatures of IMBHs in galactic nuclei. In section~\S\ref{sec:gw} we discuss gravitational wave signatures of our model potentially detectable with LIGO and LISA.
\label{sec:conclusions} If Intermediate mass black holes (IMBHs) can grow efficiently in AGN disks, the AGN host should exhibit myriad observational signatures. IMBHs that open gaps in AGN disks will exhibit strong observational parallels with gapped protoplanetary disks and may be detectable near merger in the broad FeK$\alpha$ line. LINER activity may be due to a weakly accreting MBH binary in a large disk cavity. If IMBHs do not open a gap, detection depends on signatures of accretion onto the IMBH (including tidal disruption events). We summarize observational signatures and compare them to current data where possible or suggest future observations.
14
3
1403.6433
If intermediate-mass black holes (IMBHs) grow efficiently in gas discs around supermassive black holes, their host active galactic nucleus (AGN) discs should exhibit myriad observational signatures. Gap-opening IMBHs in AGN discs can exhibit spectral features and variability analogous to gapped protoplanetary discs. A gap-opening IMBH in the innermost disc imprints ripples and oscillations on the broad Fe Kα line which may be detectable with future X-ray missions. A non-gap-opening IMBH will accrete and produce a soft X-ray excess relative to continuum emission. An IMBH on a retrograde orbit in an AGN disc will not open a gap and will generate soft X-rays from a bow-shock `headwind'. Accreting IMBH in a large cavity can generate ULX-like X-ray luminosities and LINER-like optical line ratios from local ionized gas. We propose that many LINERs house a weakly accreting MBH binary in a large central disc cavity and will be luminous sources of gravitational waves (GW). IMBHs in galactic nuclei may also be detected via intermittent observational signatures including: UV/X-ray flares due to tidal disruption events, asymmetric X-ray intensity distributions as revealed by AGN transits, quasi-periodic oscillations and underluminous Type Ia supernovae. GW emitted during IMBH inspiral and collisions may be detected with eLISA and LIGO, particularly from LINERs. We summarize observational signatures and compare to current data where possible or suggest future observations.
false
[ "future X-ray missions", "asymmetric X-ray intensity distributions", "gas discs", "AGN discs", "protoplanetary discs", "-", "rays", "ULX-like X-ray luminosities", "a soft X-ray excess relative", "UV/X-ray flares", "local ionized gas", "supermassive black holes", "myriad observational signatures", "intermittent observational signatures", "observational signatures", "future observations", "underluminous Type Ia supernovae", "quasi-periodic oscillations", "AGN transits", "continuum emission" ]
8.001899
3.539277
39
525930
[ "Fletcher, Leigh N.", "Irwin, Patrick G. J.", "Barstow, Joanna K.", "de Kok, Remco J.", "Lee, Jae-Min", "Aigrain, Suzanne" ]
2014arXiv1403.4436F
[ "Exploring the Diversity of Jupiter-Class Planets (Discussion Meeting Contribution)" ]
1
[ "-", "-", "-", "-", "-", "-" ]
[ "2021AJ....162...30S" ]
[ "astronomy" ]
16
[ "Astrophysics - Earth and Planetary Astrophysics" ]
[ "1980PThPh..64..544M", "2000imas.book.....R", "2004ApJ...611..587L", "2004jpsm.book...79W", "2005ARA&A..43..195K", "2006MNRAS.368.1087B", "2007ApJ...656.1136S", "2008JQSRT.109.1136I", "2008MNRAS.385..109P", "2009MNRAS.398..985H", "2009sfch.book..161W", "2010AIPC.1273..496L", "2010ARA&A..48..631S", "2011MNRAS.411.2199G", "2012ApJ...749...93L", "2013ApJ...763...25M" ]
[ "10.48550/arXiv.1403.4436" ]
1403
1403.4436_arXiv.txt
The high-temperature hydrogen-rich atmospheres of extrasolar giant planets (EGPs) in close orbits around their parent stars have made them ideal candidates for preliminary spectroscopic characterisation. Giant exoplanets (i.e., Neptune sized and larger) appear to be commonplace: from a catalog of confirmed exoplanets with mass determinations\footnote{http://exoplanet.eu}, 75\% of all planets discovered to date have masses larger than Saturn (0.3$M_J$), 53\% are more massive than Jupiter, and 67\% are within 1 AU of their parent star. However, this could simply be the result of observational bias for the larger giants - when the list of Kepler candidate objects is included \citep{13fressin, 13batalha}, Neptune-sized giants could comprise a significant percentage \citep[30\% or larger,][]{13fressin} of planetary objects beyond our solar system, restricting Jupiter-sized objects to 5\% or less. This article will focus on the taxonomy of hydrogen-rich gaseous exoplanets, in particular the Jupiter-class with radii $>6R_E$ \citep[following the prescription of the Kepler team,][]{13fressin}, to which our own Jupiter (11 $R_E$) and Saturn (9.1 $R_E$) belong. However, many of the conclusions may be equally valid for Neptune-class objects ($2-6R_E$) or hydrogen-rich super earths (1.25-2 $R_E$) that may dominate the planetary populations beyond our solar system. The conditions revealed on these worlds are significantly different from the gas giants of our own solar system, and yet these hot ($\approx2500$ K) and cold ($\approx100$ K) jovians must exist on a continuum of planetary types that can be categorised in terms of a range of atmospheric phenomena. Transit spectroscopy of a handful of EGPs tentatively revealed the presence of simple molecules of hydrogen, carbon and oxygen \citep[water,methane, CO and CO$_2$,][]{07tinetti,08swain,09swain, 11gibson}, along with sodium \citep{02charbonneau}, atomic H and other neutral species in their upper atmospheres. Some of these conclusions have been refined and questioned in subsequent years, and observers have gone to great lengths to confirm or refute these molecular detections to move the field forward. Atmospheric models have revealed the importance of strong stellar insolation, the possible presence of atmospheric hazes \citep[e.g.,][]{08pont}, stratospheric thermal inversions \citep{09swain}, atmospheric winds \citep{10snellen} and longitudinal temperature contrasts \citep{07knutson}. Approximately 25-30\% of the confirmed planets are known to transit their host stars, permitting spectroscopic characterisation via the transit method biased towards highly-irradiated, close-in giant planets on short period orbits. A considerably smaller number of directly-imaged planets well separated from their host stars on longer orbits (young and hot worlds) are presently available for spectroscopic characterisation, but it is hoped that this number could increase in the near future. So, given the sparse information we have on the composition, dynamics and chemistry of these EGPs, development of classification systems may seem premature. But they serve a useful purpose, allowing us to bring a degree of order to the patterns in the emerging pantheon of planetary types being discovered today, providing a vernacular for their discussion \citep[by analogy to the categorisation of brown dwarfs,][]{05kirkpatrick} and testable hypotheses to be addressed by future spectroscopic missions. A series of one- and two-dimensional classification systems have been proposed in the decade since the work of \citet{03sudarsky}, who used a solar composition and thermochemical equilibrium to study the influence of stellar irradiation and its implications for atmospheric composition and cloud formation. Put simply, the higher the irradiation, the hotter the atmosphere and the more refractory species are released from the condensed phase to interact with the emission from the planetary photosphere (broadly speaking, the middle atmosphere of the planet from the stratosphere to the upper troposphere). At the highest temperatures, oxides of titanium and vanadium would be available \citep[if the planet is oxygen-rich,][]{12madhu} to serve as strong UV/visible absorbers, generating stratospheric temperature inversions \citep[e.g.,][]{03hubeny} and leading \citet{08fortney} to develop a one-dimensional classification system (pM and pL classes, by analogy to M and L brown dwarfs) based on the presence or absence of a thermal inversion. As we shall see in Section \ref{spectra}, existing observations are generally not robust enough to confirm the stratospheric temperature gradients, but several studies have identified highly-irradiated EGPs with no thermal inversions at all\citep{11madhu}, betraying the simplicity of any irradiation-driven categorisation scheme. Indeed, Section \ref{processes} reveals several other factors influencing the atmospheric temperature and composition, and good reviews of these processes can be found in \citet{10seager} and \citet{10burrows}. Gravitational settling and cold-trapping (the same process that keeps the Earth's stratosphere free of water vapour) may deplete TiO and VO from EGP atmospheres and limit their role in forming stratospheric inversions \citep{09spiegel}. Furthermore, an oxygen-poor, carbon-rich atmosphere would limit the production of these oxides, as well as having dramatic effects on other carbon species \citep{09helling, 11madhu, 12madhu}. Chromospheric activity of the parent stars (e.g., flares and cosmic rays) can either inhibit photochemistry and destroy potential UV/visible absorbers, or excite additional ion chemistry due to excess charged particle bombardment, which can generate upper atmospheric hazed that contribute to the radiative budget \citep{09zahnle_soot, 10knutson, 13moses}, although the residency times and importance of these species for generating thermal inversions is uncertain. Other potential stratospheric absorbers may be tied up in condensed phases not yet considered in models. Different thermal structures will affect an atmosphere's susceptibility to photochemistry and vertical mixing \citep{10line, 11moses, 13moses}. Vertical and horizontal mixing by eddy diffusion, convection and wave propagation would also be responsible for moving energy and material from place to place, causing the composition to deviate from the expectations of equilibrium \citep[e.g.,][]{10line,11moses,10showman}. Each of these processes could provide additional dimensions to a classification scheme for EGPs. Of particular note is the recent two-dimensional scheme devised by \citet{12madhu} and \citet{13moses}, which uses both the stellar irradiance and the chemical dependence on the C/O ratio, described below. Although the true atmospheric temperature-pressure ($T(p)$) profile is desirable for a complete discussion of chemistry and dynamics, in reality we must use a proxy, the equilibrium temperature $T_{eq}$ based on the stellar luminosity, $L$; the bond albedo $a$; the degree of horizontal temperature redistribution $f$; the orbital distance $r$ and Stefan-Boltzmann constant, $\sigma$ \citep{03sudarsky}: \begin{equation} T_{eq}=\left[ \frac{L(1-a)}{16\pi\sigma r^2f}\right]^{1/4} \end{equation} In Fig. \ref{Teqm}, we estimate $T_{eq}$ for all planets discovered to date, accounting for the uncertainty in the bond albedo (from zero to a Solar-System-like value of 30\%) and the efficiency of redistribution ($f=1$ for a full redistribution of incoming flux; $f=0.5$ if radiation is only emitted from the dayside atmosphere). This differs from the effective temperature ($T_{eff}$) which also takes into account an internal boundary flux that varies according to the age and thermal history of the planet, such as the excess luminosity of the giant planets in our own solar system \citep{91pearl}. The number of planets discovered in a particular $T_{eq}$ category is also presented in Fig. \ref{Teqm}. Although a poor proxy for the true $T(p)$, this quantity does at least permit preliminary categorisations \citep[e.g.,][]{03sudarsky,08fortney,12madhu}, and will be used as a guide in the text that follows. \begin{figure*}[tbp] \centering \includegraphics[width=16cm]{exo_Teqm.pdf} \caption{The diversity of Jupiter-class exoplanets discovered to date with constrained masses (all data from the catalogue hosted at exoplanet.eu). Panel (a) plots the equilibrium temperature against planetary mass (with Jupiter as a red circle and Saturn as a blue circle for comparison); (b) shows the range of equilibrium temperatures discovered for each stellar type from F to M; and (c) shows the number of planets discovered to be occupying a particular equilibrium temperature range. Although $T_{eq}$ is a poor proxy for the true atmospheric temperatures, it does permit a first-order attempt at classifying these EGPs. Error bars depict the range of $T_{eq}$ for bond albedoes between 0.0 and 0.3, and the difference between purely dayside re-emission for full energy redistribution). Panel (d) shows the planetary mass and orbit of all exoplanets discovered to date. } \label{Teqm} \end{figure*}
This review has demonstrated that the Jupiter-class of EGP represents a broad continuum of planetary types, with bulk composition, equilibrium and disequilibrium chemistry, and the complex sequence of condensates all competing to shape the emergent spectra. Classification schemes involving stellar irradiance and condensate formation, bulk composition and chemistry, and the influence of strong stellar activity have all been proposed \citep{03sudarsky, 08fortney, 10knutson, 12madhu, 13moses} and capture the essence of the EGP classification problem. All these schemes are model-dependent, and they are all hampered by the lack of observational constraints, particularly for the cooler jovians bridging the gap between our Solar System and the two `Rosetta Stones', HD 189733b ($T_{eq}=1200_{-105}^{+230}$ K, a hazy Jupiter lacking a thermal inversion, orbiting a cool chromospherically-active K star) and HD 209458b ($T_{eq}=1440_{-125}^{+270}$, a cloud-free, hot metallic Jupiter with a thermal inversion orbiting a G star). Classification requires a robust ensemble of atmospheric compositional types, under varying irradiation conditions (i.e., different stellar types, power and orbital radii). It also requires spectroscopy that is both accurate and sufficiently broadband to (i) determine the continuum formed by atmospheric temperatures and hazes and (ii) unambiguously detect the presence of molecular species. Space-borne spectroscopy from the UV to the infrared, potentially from JWST or ECHO \citep{12tinetti}, could begin to validate some of the early findings on irradiated EGPs and extend coverage to planets with longer orbital periods and smaller stellar influences (see Fig. \ref{exocloud} for the range of $T_{eq}$ within reach of such transit studies). Based on the current statistics of targets suitable for EChO, up to 50\% of the mission time could be devoted to the Jupiter class \citep{12tinetti}. Hot metallic Jupiters and silicate cloud Jupiters will be well-sampled across all stellar types (the \textit{hot} sample, $T_{eq}>1800$ K), but sulphide cloud Jupiters, cloud free Jupiters with strong alkali lines ($700<T_{eq}<1800$, the \textit{temperate} sample) and water-cloud jovians (the \textit{cool} sample) could also be targeted. Cooler Jupiters must be observed around cooler M and K stars, so that their orbits are still sufficiently short to permit the observation of multiple transits (Fig. \ref{exocloud}). EGP categorisation schemes will allow us to make optimal selections of targets for such a mission, and only by assembling a reference collection of well-characterised EGP photospheres for a range of different stellar flux conditions can we begin to test the predictions of the EGP schemes reviewed here. We may find that certain planets could be considered archetypes of their classes, in the same way as Jupiter is seen for our solar system giant planets. Ultimately it seems certain that understanding this `continuum of jovians' will lead us to view our own cold giant planets in an entirely new way.
14
3
1403.4436
Royal Society Discussion Meeting (2013) `Characterizing exoplanets'. Of the 900+ confirmed exoplanets discovered since 1995 for which we have constraints on their mass (i.e., not including Kepler candidates), 75% have masses larger than Saturn (0.3MJ), 53% are more massive than Jupiter, and 67% are within 1 AU of their host stars. And yet the term `hot Jupiter' fails to account for the incredible diversity of this class of object, which exists on a continuum of giant planets from the cool jovians of our own solar system to the highly-irradiated, tidally-locked hot roasters. We review theoretical expectations for the temperatures, molecular composition and cloud properties of Jupiter-class objects under a variety of different conditions. We discuss the classification schemes for these Jupiter-class planets proposed to date, including the implications for our own Solar System giant planets and the pitfalls associated with classification at this early stage of exoplanetary spectroscopy. We discuss the range of planetary types described by previous authors, accounting for: (i) thermochemical equilibrium expectations for cloud condensation and favoured chemical stability fields; (ii) the metallicity and formation mechanism for these giant planets; (iii) the importance of optical absorbers for energy partitioning and the generation of a temperature inversion; (iv) the favoured photochemical pathways and expectations for minor species (e.g., saturated hydrocarbons and nitriles); (v) the unexpected presence of molecules due to vertical mixing of species above their quench levels; and (vi) methods for energy and material redistribution throughout the atmosphere (e.g., away from the highly irradiated daysides of close-in giants). Finally, we will discuss the benefits and flaws of retrieval techniques for establishing a family of atmospheric solutions that reproduce the available data.
false
[ "giant planets", "favoured chemical stability fields", "energy partitioning", "thermochemical equilibrium expectations", "minor species", "exoplanetary spectroscopy", "our own Solar System giant planets", "Jupiter", "different conditions", "vertical mixing", "cloud properties", "cloud condensation", "species", "theoretical expectations", "expectations", "optical absorbers", "Kepler candidates", "Solar System", "Jupiter-class objects", "object" ]
7.769171
14.362455
109
437480
[ "Reynoso, M. M." ]
2014A&A...564A..74R
[ "A two-zone approach to neutrino production in gamma-ray bursts" ]
10
[ "Department of PhysicsUniversity of Athens, Panepistimiopolis, 15783, Zografou, Greece ; Instituto de Investigaciones Físicas de Mar del Plata (CONICET - UNMdP), Facultad de Ciencias Exactas y Naturales, Universidad Nacional de Mar del Plata, Dean Funes 3350, (7600), Mar del Plata, Argentina" ]
[ "2014A&A...569A..58W", "2014MNRAS.442.3026P", "2014MNRAS.444.2186P", "2014MNRAS.445..570P", "2014PhRvD..90j3003W", "2015MNRAS.448.2412P", "2015PhRvD..92h5047K", "2016APh....82...10R", "2020A&A...641A..29G", "2022NatRP...4..697G" ]
[ "astronomy" ]
4
[ "radiation mechanisms: non-thermal", "neutrinos", "gamma-ray burst: general", "Astrophysics - High Energy Astrophysical Phenomena" ]
[ "1962SvA.....6..317K", "1968PhRv..167.1159J", "1970RvMP...42..237B", "1983Afz....19..323A", "1990ApJ...362...38B", "1994ApJ...430L..93R", "1994PhRvD..50.5710F", "1996ApJ...473..998F", "1997ApJ...490...92K", "1997PhRvL..78.2292W", "1998A&A...333..452K", "1999A&A...347..370D", "1999APh....10..185P", "2000ApJ...530..292M", "2000ApJ...541..707W", "2000CoPhC.124..290M", "2001ApJ...551..249D", "2003ApJ...586...79A", "2004APh....20..429G", "2004PhRvL..93r1101R", "2004RvMP...76.1143P", "2005PhRvL..95f1103A", "2006A&A...457..763G", "2006PhRvD..73f3002M", "2006PhRvD..74c4018K", "2006RPPh...69.2259M", "2007A&A...462..173M", "2007PhRvD..75l3005L", "2007PhRvD..76l3001M", "2009A&A...493....1R", "2010ApJ...725L.121A", "2011A&A...531A..30R", "2012A&A...542A.125B", "2012APh....35..508B", "2012ApJ...746..164M", "2012ApJ...752...29H", "2012JCAP...11..058G", "2012JHEP...12..123G", "2012MNRAS.422.2474D", "2013A&A...558A.142V", "2013ApJ...768..197I", "2014PhRvD..89h3004L" ]
[ "10.1051/0004-6361/201323159", "10.48550/arXiv.1403.3020" ]
1403
1403.3020_arXiv.txt
Gamma-ray bursts (GRB) are intense and brief flashes of gamma rays that last from a fraction of a second to tens of seconds, releasing energies as high as $10^{51-53}${\rm erg} \cite{piran2004,meszaros2006}. {While short bursts with durations $t_{\rm GRB}^{\rm obs}\sim 2$ s are believed to be caused by the merger of compact stars in a binary system, long bursts ($t_{\rm GRB}^{\rm obs}\gtrsim 10$ s) are thought to be triggered by the collapse of a massive star into a black hole.} In the most accepted scenario, the prompt emission corresponding to the observed burst is supposed to come from synchrotron and/or inverse Compton emission of electrons that are accelerated in internal shocks of ejecta with various Lorentz factors, $\Gamma\sim 100-1000$ \cite[e.g.][]{rees1994,fenimore1996,kobayashi1997}. However, this is not the only possibility: photospheric models and acceleration by reconnection have also been proposed to explain such emission \cite[e.g.][]{meszaros2000,giannios2006,gao2012}.% Neutrino production in GRBs is expected if, for instance, protons are co-accelerated with the electrons responsible for the prompt emission. Then, $p\gamma$ and $pp$ interactions in the the baryon rich flow would lead to pion production, and thus to neutrinos \citep[e.g.][]{waxmanbahcall1997,guetta2004,murase2006}. In more recent studies, new calculations have been developed to obtain the possible neutrino flux under different assumptions % \cite[e.g.][]{hummer2012,murase2012,baerwald2012,he2012}. It has also been proposed that in an earlier stage, while the jet is still propagating inside the collapsing star or just outside its surface, shocks may develop but without an observable photon counterpart, and only neutrinos would escape \cite{razzaque2004,ando2005,vieyro2013}. Another well studied possibility is the generation of neutrinos during the afterglow phase \cite[e.g.][]{waxman2000,dai2001}, which corresponds to a delayed low energy emission that occurs from hours to days after the prompt emission, and is commonly explained by external shocks with the interstellar medium. In the present work, we focus on the production of prompt neutrinos, considering the effects of a generic acceleration process acting on all charged particles, including the secondary pions and muons. To do this, we adopt a simple model with two zones: an acceleration zone and a cooling one, and we assume that the particles escaping from the former are injected into the latter. We find that in the cases where the escape rate is slower than the acceleration rate, then the synchrotron emission from the electrons in the acceleration zone can yield a flux that is consistent with GRB observations. Then, the co-accelerated protons can produce significant amounts of pions by $pp$ and $p\gamma$ interactions depending on the power injected in protons. The decaying muons can undergo acceleration in the cases of higher magnetic fields, for which their acceleration rate becomes higher than their decay rate. For illustration, we compute the diffuse neutrino flux that would be expected from GRBs under some different assumptions on the Lorentz factor and on the escape rate, and we compare these results with the the Waxman-Bahcall GRB flux \cite{waxmanbahcall1997} and with the data of the recent neutrino detection by IceCube \cite{aartsen2013}. This work is organized as follows. In Section 2, we describe the basic assumptions of the model, and in Section 3 we show the results obtained for the particle distributions: protons, electrons, pions and muons. In Section 4 we show illustrative results for the predicted broadband GRB photon flux, and in Section 5 we compute the corresponding diffuse fluxes of prompt neutrinos. The final comments are made in Section 6.
We have implemented a simple two-zone model in order to study the generation of high energy neutrinos associated with the prompt GRB emission. Using standard values for the magnetic field and size of the emission region, our model can account for the possible effect of the acceleration of secondary particles. In particular, we found that muons can efficiently gain energy if the magnetic field is strong enough, but still within attainable values in the context of GRBs. We note that these effects cannot be described with previous one-zone models that deal with neutrino emission in a magnetized environment \cite[e.g.][]{magnetic2009,baerwald2012}, in which the acceleration rate is only used to fix the maximum energy of the primary electrons and protons. As recognized in previous works \cite[e.g.][]{kirk1998}, particle acceleration can be accounted for using two zones and assuming that particles can escape from the acceleration zone to the cooling zone. {We have not considered that particles in the cooling zone can further escape to a third zone in order not to miss their photon and neutrino output.} {The present model also differs from previous two-zone models in that the size of both zones are equal, and with a value derived from variability considerations. Including adiabatic losses for protons provides a mechanism for their faster cooling, on a timescale similar to the dynamical time, e.g. the one associated with the duration of the shell collision event in the internal shock scenario. A variation of the present model could be implemented by including a convective term in the kinetic equation for the cooling zone \cite[e.g.][]{leptoh2011}. This would prevent us from having to impose a fixed size for the cooling zone, since particles of different species and energies would reach different distances as they cool.} In the context studied here, we have found that if the escape rate is less than the acceleration rate ($\xi_{\rm esc}<1$), then the synchrotron emission from electrons in the acceleration zone dominates and can be the responsible for the usual GRB emission. Otherwise, for faster escape rates, the synchrotron emission from electrons in the cooling zone would dominate but with a spectrum too wide, which would greatly exceed the typical GRB emission at lower energies. As can be seen in Fig. \ref{fig6:E2dfluga_g100_g300}, for lower values of $\xi_{\rm esc}$, we obtain less significant electron synchrotron components from the cooling zone, and the bump corresponding to the acceleration zone falls within the correct energy range, as compared with the broken power-law benchmark. By varying the acceleration efficiency of the different injection events (such as shell collisions in the internal shock model), different maximum energies for the electrons could be achieved, and their synchrotron emission would cover a window in the gamma-ray spectrum to be consistent with a full burst. In such cases with a low escape rate, we found that a neutrino component arising from the acceleration zone mainly by $p\gamma$ interactions becomes dominant at the highest neutrino energies, which in the examples shown reached $\sim 10^6$ GeV and can account for the recent IceCube data. Some tasks could help make a more accurate calculation of the diffuse neutrino background in the context of the present type of models for GRBs: try to reproduce the observed gamma-ray spectrum of particular bursts by adjusting the number of acceleration events (peaks in the lightcurve) and the acceleration efficiency, and also to consider the probability of occurrence of bursts with different Lorentz factors. We leave these points for future work, along with the possible application of the model to other type of astrophysical sources.
14
3
1403.3020
Context. Gamma-ray bursts (GRB) are the most powerful events in the universe. They are capable of accelerating particles to very high energies, so are strong candidates as sources of detectable astrophysical neutrinos. <BR /> Aims: We study the effects of particle acceleration and escape by implementing a two-zone model in order to assess the production of high-energy neutrinos in GRBs associated with their prompt emission. <BR /> Methods: Both primary relativistic electrons and protons are injected in a zone where an acceleration mechanism operates and dominates over the losses. The escaping particles are re-injected in a cooling zone that propagates downstream. The synchrotron photons emitted by the accelerated electrons are taken as targets for pγ interactions, which generate pions along with the pp collisions with cold protons in the flow. The distribution of these secondary pions and the decaying muons are also computed in both zones, from which the neutrino output is obtained. <BR /> Results: We find that for escape rates lower than the acceleration rate, the synchrotron emission from electrons in the acceleration zone can account for the GRB emission, and the production of neutrinos via pγ interactions in this zone becomes dominant for E<SUB>ν</SUB> &gt; 10<SUP>5</SUP> GeV. For illustration, we compute the corresponding diffuse neutrino flux under different assumptions and show that it can reach the level of the signal recently detected by IceCube.
false
[ "detectable astrophysical neutrinos", "neutrinos", "pγ interactions", "particle acceleration", "IceCube", "cold protons", "different assumptions", "high-energy neutrinos", "electrons", "protons", "the corresponding diffuse neutrino flux", "the acceleration zone", "gt", "strong candidates", "the neutrino output", "pions", "a cooling zone", "GRB", "their prompt emission", "escape rates" ]
5.650428
0.147782
13
482667
[ "Kholikov, S.", "Serebryanskiy, A.", "Jackiewicz, J." ]
2014ApJ...784..145K
[ "Meridional Flow in the Solar Convection Zone. I. Measurements from GONG Data" ]
46
[ "National Solar Observatories, Tucson, AZ 85719, USA", "Ulugh Beg Astronomical Institute, Uzbek Academy of Science, Tashkent 100052, Uzbekistan", "Department of Astronomy, New Mexico State University, Las Cruces, NM 88003, USA" ]
[ "2014A&A...570A..43K", "2014GeoRL..41.5361D", "2015A&A...576A..26K", "2015ApJ...803L..17G", "2015ApJ...805..133J", "2015ApJ...805..165L", "2015ApJ...806..169B", "2015MNRAS.451.1528P", "2015SoPh..290..659K", "2015SoPh..290.3113K", "2016AdSpR..58.1457Z", "2016AdSpR..58.1490P", "2016AdSpR..58.1589D", "2016ApJ...824...49B", "2016ApJ...824..128G", "2016ApJ...828...91D", "2016IAUTA..29..278C", "2016SoPh..291..731Z", "2017A&A...601A..46L", "2017ApJ...835....9B", "2017ApJ...838...53B", "2017ApJ...845....2B", "2017ApJ...849..144C", "2017ApJ...849..160H", "2018A&A...619A..99L", "2018Ap&SS.363..266L", "2018ApJ...853..161C", "2018ApJ...854...67P", "2018ApJ...854...85W", "2018ApJ...860...48L", "2018IAUS..335..183H", "2018SoPh..293..134B", "2018arXiv180910776K", "2019ASSL..458.....A", "2019ApJ...871L..20K", "2019Entrp..21..140W", "2021ApJ...909...66H", "2021ApJ...911...54B", "2021ApJ...911...90S", "2021PNAS..11822518V", "2022ApJ...933..109Z", "2022ApJ...934..161S", "2022tess.conf11009Z", "2023IAUS..362..333K", "2023arXiv230913076S", "2023arXiv231211699T" ]
[ "astronomy" ]
2
[ "Sun: helioseismology", "Sun: interior", "Sun: oscillations", "Astrophysics - Solar and Stellar Astrophysics" ]
[ "1979SoPh...63....3D", "1982ApJ...256..316G", "1989Sci...245..712W", "1991ApJ...383..431W", "1991sia..book..479L", "1993Natur.362..430D", "1993SoPh..147..207K", "1995A&A...303L..29C", "1996ApJ...460.1027H", "1997ASSL..225..241K", "1997Natur.390...52G", "1998ApJ...508L.105B", "1999ApJ...510L.153G", "1999ApJ...512..458B", "1999ApJ...518..508D", "2000PhDT.........9G", "2001ApJ...559L.175C", "2002ApJ...570..855H", "2002ApJ...575L..47B", "2002ApJ...577L..53W", "2003ApJ...585..553B", "2004ApJ...603..776Z", "2005ApJ...630.1206C", "2006ApJ...638..576G", "2006SoPh..236..227Z", "2008SoPh..251..241G", "2008SoPh..252..235G", "2009ASPC..416..103D", "2010ApJ...713L..16G", "2010ApJ...717..488B", "2010ApJ...720.1030C", "2010ApJ...725..658U", "2010Sci...327.1350H", "2011Natur.471...80N", "2012ApJ...746...65D", "2012ApJ...749L...5Z", "2012ApJ...749L..13H", "2012ApJ...760L...1B", "2012ApJ...761L..14R", "2012SoPh..275..375Z", "2013ApJ...762..132H", "2013ApJ...774L..29Z", "2013ApJ...778L..38S", "2014SoPh..289.1077K" ]
[ "10.1088/0004-637X/784/2/145", "10.48550/arXiv.1403.5026" ]
1403
1403.5026_arXiv.txt
Meridional circulation plays a critical role in models of solar dynamo, magnetic flux transport, and the solar cycle \citep{glatzmaier1982,wang1989,wang1991,choudhuri1995,dikpati1999,wang2002,nandy2011}. It is well established observationally that meridional flow is poleward in each hemisphere with an amplitude of about $10-20~{\rm m\,s^{-1}}$ in the near-surface layers, peaking in strength at mid latitudes \citep{duvall1979,hathaway1996,braun1998,hernandez1999,basu1999,hernandez2006,basu2010,hathaway2010,ulrich2010}. Since mass does not pile up at the poles, it is believed that a return equatorward flow in both hemispheres is operating somewhere in the convection zone, likely near its base. One of the most promising and complete attempts to measure this meridional circulation was during the graduate work of P.~Giles \citep{giles1997,giles2000}. Using the SOHO spacecraft's Michelson Doppler Imager (MDI) helioseismic data, Giles found that the poleward meridional flows continued throughout almost all of the convection zone and that there was indirect evidence of a return equatorward flow near the tachocline of a few ${\rm m\,s^{-1}}$. His methods and analysis imposed a constraint of mass conservation. Thus, the picture that emerged was of two closed circulating flows, one cell in each hemisphere, that diverge from the equator at the surface and converge toward the equator in the deep interior. Since then, other helioseismology studies using a variety of techniques have offered many differing views. For example, \citet{chou2001,beck2002,chou2005} observe an additional ``cell'' of meridional circulation at mid latitudes near the location of the active sunspot latitudes, which is divergent and varies in time. Also, \citet{zhao2004,hernandez2010} found that such a cell has a convergent flow field \citep{cameron2010}. Indeed, large-scale flow profiles (in both meridional and zonal directions) have been found to vary rather strongly with the solar cycle, and several studies have found that the amplitude of the flow is anti-correlated with the strength of the cycle \citep[e.g.,][]{komm1993,chou2001,haber2002,basu2003,hernandez2008,gizon2008}. The latitudinal extent of the surface poleward flow has widely varied in the two previous cycles, and some helioseismic measurements indicate a high-latitude, reverse, \emph{equatorward} surface component \citep{dikpati2012}. To add to the complexity, recent observations have shown an increasing polar flow magnitude as one probes deep into the convection zone \citep{kholikov2012}, and \citet{hathaway2012} place the equatorward return flow at a depth of 70~Mm. Recently \citet{zhao2012b} observed a new systematic center-to-limb signal in time-distance measurements \citep{duvall1993}, which may play a key role in obtaining reliable deep meridional flow measurements and be one of the sources of the discrepant results mentioned above. The approach of \citet{zhao2012b} was to remove the systematic travel-time shifts found in the east-west measurements, after rotation is removed, from the meridional (north-south) measurements. This correction led to consistent helioseismic measurements using several different observables. While the source of this signal is not completely understood, it could be related to existing observational limitations like changes of the line formation heights across the solar disk, which produce additional acoustic travel-time shifts in cross-correlation measurements between different locations. \citet{baldner2012} showed that the effect of the vertical flows from convection in the outer solar convection zone can similarly affect travel-time measurements. Subsequently, \citet{zhao2013} applied their local techniques to measure two meridional circulation cells in the solar convection zone, while \citet{schad2013} implemented a new global helioseismic analysis that resulted in evidence of a complex multicellular velocity structure. These new and exciting findings from space-based data present a potentially revised view of these important large-scale flows. This paper is the first in a series where we explore meridional circulation using time-distance helioseismology applied to Global Oscillation Network Group (GONG) data. Here we describe in detail the travel-time measurement procedure we implement, which is non-standard and differs from the methods of \citet{zhao2012,hartlep2013}, for example. We use more than 600 daily sets of GONG velocity images to probe deep into the convection zone. In order to decrease possible geometric and observational artifacts we have selected dates with a duty cycle of more than 85\% and time periods when the solar tilt angle $B_0 \le 4^\circ$. These strict requirements substantially decrease the amount of data that can be used. We assume that the center-to-limb systematic mentioned above is the same in any direction on the solar disk and we compute it only using the equatorial region of the observations. Travel-time differences are computed for north-south flows and corrected by subtracting the east-west signal. We find strong evidence of a change of sign in the travel-time differences at mid latitudes and depths of about 200~Mm beneath the surface, and compelling evidence that travel times may change sign (thus signaling a flow reversal) at shallower depths of about 50-60~Mm. In Section~\ref{sec:data} we describe the data and analysis procedure, with results and discussion provided in Section~\ref{sec:res}.
\label{sec:res} \begin{figure} \centerline{ \includegraphics[width=.5\textwidth,clip=]{fig21} \includegraphics[width=.5\textwidth,clip=]{fig22}} \vspace{-.4\textheight} \centerline{\hspace{.06\textwidth}\bf\Large\color{black}{(a)}\hspace{.45\textwidth}\bf\color{black}{(b)}\hfill} \vspace{.35\textheight} \caption{Travel-time difference maps obtained using 652 daily sets of Doppler velocity images. Column (a) shows the SN, EW, and SN-EW contour maps from top to bottom, respectively. Column (b) plots the corresponding measurement uncertainties associated with each panel in column (a). Note the $x$ axis in the middle panel in each column is the longitude, with the same numerical scale values as shown for latitude ($\pm 75\arcdeg$). Hatched regions show where no measurements were computed due to limb constraints.\label{maps}} \end{figure} \begin{figure} \centerline{ \includegraphics[width=.5\textwidth,clip=]{fig31} \includegraphics[width=.5\textwidth,clip=]{fig32}} \vspace{-.2\textheight} \centerline{\hspace{.07\textwidth}\bf\large\color{black}{(a)}\hspace{.47\textwidth}\bf\color{black}{(b)}\hfill} \vspace{.17\textheight} \centerline{ \includegraphics[width=.5\textwidth,clip=]{fig41} \includegraphics[width=.5\textwidth,clip=]{fig42}} \vspace{-.1\textheight} \centerline{\hspace{.07\textwidth}\bf\large\color{black}{(c)}\hspace{.47\textwidth}\bf\color{black}{(d)}\hfill} \vspace{.05\textheight} \caption{Cuts through depth and latitude of corrected SN travel-time differences. The top row panels (a) and (b) show the travel-time differences in each hemisphere as a function of measurement distance ($\Delta$) for the latitude range averaged over the $10^\circ$ band noted in the figure. A proxy for the lower turning point depth for each travel distance is shown on the upper $x$ axis. Highlighted in gray are the travel distances shown in the corresponding plots below. Panels (c) and (d) shows travel times as a function of latitude averaged over an interval in distances of $5^\circ$. The latitude ranges in panels (a) and (b) are given in the gray boxes of panels (c) and (d). The uncertainties are shown for all cases and are plotted only at staggered data points for clarity.\label{cuts}} \end{figure} The top left panel of Fig.~\ref{maps} shows an average over 652 days of SN travel-time differences presented as a function of latitude and travel distance. Each point at a given travel distance corresponds to the middle position between a point and an arc in our cross-correlation scheme. To avoid very high latitude information where the endpoints of the cross correlations lie, the measurements are cut off as a function of distance. The uncertainties are given in the second column, computed from the dispersion in individual measurements for each longitude and each day. These are typically a very small percentage of the averaged signal. Signatures of poleward meridional flow in each hemisphere are clearly seen in Fig.~\ref{maps}(a). The color convention in this figure is such that blue is consistent with a flow toward the North Pole, and red a flow toward the South Pole. Indeed, in addition to a peak at mid latitudes as expected, an increase in the travel-time difference with depth (i.e., travel distance) is also observed. We expect this to be due to one or several systematics. To explore this further, EW travel times computed from the same dataset are shown in the middle panel of Fig.~\ref{maps} as a function of longitude on the $x$ axis. The EW map has been symmetrized about the central meridian, as we expect there to be no significant differences between the two (east/west) hemispheres since the data have been tracked to account for differential rotation. These measurements show a similar pattern of center-to-limb variation as the SN map. \citet{zhao2012} reported a very detailed analysis of travel-time measurements from different observables. Since they found that the shape and magnitude of center-to-limb variations is quite different for Doppler, continuum, line core and line depth of HMI measurements, one might conclude that these variations are not caused by any large-scale sub-surface flow of solar origin. Here we follow the same procedure and ``correct'' the SN measurements by subtraction of the EW measurements, the result shown in the bottom panel of Fig.~\ref{maps}(a). This correction removes the tendency of the travel times to increase with depth. Furthermore, some evidence of sign changes can be seen. Figure~\ref{cuts} shows various cuts through the travel-time difference maps. Panels (a) and (c) are cuts at lower latitudes and shorter travel distances, while panels (b) and (d) are for mid latitudes and larger travel distances. These figures confirm that travel-time differences are strongest at mid latitudes around $30^\circ$ for a range of depths, as has been observed in past studies. This representation shows a clear yet peculiar asymmetry between the northern and southern hemispheres. Most importantly, we also observe evidence that a change in sign occurs in the measurements for two cases: (1) at high latitudes in each hemisphere for travel distances greater than about $15^\circ$; and (2) for large distances for most latitudes greater than about $20^\circ$ in each hemisphere. Indeed, if large-scale flows are responsible for these signals, Figs.~\ref{cuts}(a)-(b) show a tendency for the flow to approach a change of sign at skip distances of $15^\circ-20^\circ$ for a broad latitude range. At larger distances this signal then resurrects its poleward sense, eventually reversing again at the deepest probe depths. This very broadly suggests a multicellular structure as discussed in \citet{zhao2012} and \citet{zhao2013}, who found poleward flows down to $0.91~R_\odot$, equatorward flows in the $0.82 - 0.91~R_\odot$ range, and then poleward again beneath that. Very recent work by \citet{schad2013} reports yet another measurement of multicellular structure of the meridional flow using a different, global approach. We caution that the change in sign at all distances at the maximal latitudes considered here (most evident in Fig.~\ref{maps}) could be due to a systematic caused by the solar $B_0$ variation, as demonstrated recently in \citet{kholikov2013}. However, in the measurements here such an artifact is somewhat puzzling since we have restricted the data coverage to epochs when this angle is small. Another possible cause could simply be the use of the ad hoc correction method and any of its inherent systematics. Also evident in the measurements is a north-south hemispheric asymmetry from GONG and space-based data that has been noted in previous works \citep[e.g.,][]{zaatri2006,rightmire2012}. The real origin of the center-to-limb variation across the solar disk is not well understood at present. The work of \citet{baldner2012} in explaining it is promising. Anomalous artifacts were even identified as early as \citet{duvall2009}, who consider effects due to the finite speed of light in meridional-flow measurements as causing an overall inflow towards the disk center. While only several seconds, the proper correction for this effect individually actually tends to add to the already unphysical increasing travel-time difference signal with depth as observed in the top panel of Fig.~\ref{maps}(a). Presumably this systematic is already accounted for in the east-west subtraction correction implemented here, although more confidence in such an approach is certainly needed and is the focus of current work. Nevertheless, we find strong evidence of variations in depth of the large-scale flow in the solar convection zone. The deepest measurements where the ray path is horizontal and less sensitive to surface flows show strong evidence of a change of sign. The tendency for the travel-time shifts to approach zero at mid-convection zone depths is also intriguing, as one must recall that these measurements are integrated over depth and smoothed to some degree. Inversions may separate the two directional components of the flow and provide amplitudes and a more accurate depth structure. We have shown robust travel-time measurements of the meridional flow signature in the solar convection zone using GONG data and independent measurement techniques. Preliminary evidence of a change of sign, indicating an equatorward return flow at one or several depths in the convection zone is observed, and is approximately consistent with results found in other recently published work. Overall the findings might suggest multicellular structure in the large-scale flows in the Sun. Only a consistent inversion procedure and a very careful treatment of the systematics can unravel the significance of these trends in the measurements. A forthcoming paper will show such inversions and discuss the implications for convection-zone dynamics.
14
3
1403.5026
Large-scale plasma flows in the Sun's convection zone likely play a major role in solar dynamics on decadal timescales. In particular, quantifying meridional motions is a critical ingredient for understanding the solar cycle and the transport of magnetic flux. Because the signal of such features can be quite small in deep solar layers and be buried in systematics or noise, the true meridional velocity profile has remained elusive. We perform time-distance helioseismology measurements on several years worth of Global Oscillation Network Group Doppler data. A spherical harmonic decomposition technique is applied to a subset of acoustic modes to measure travel-time differences to try to obtain signatures of meridional flows throughout the solar convection zone. Center-to-limb systematics are taken into account in an intuitive yet ad hoc manner. Travel-time differences near the surface that are consistent with a poleward flow in each hemisphere and are similar to previous work are measured. Additionally, measurements in deep layers near the base of the convection zone suggest a possible equatorward flow, as well as partial evidence of a sign change in the travel-time differences at mid-convection zone depths. This analysis on an independent data set using different measurement techniques strengthens recent conclusions that the convection zone may have multiple "cells" of meridional flow. The results may challenge the common understanding of one large conveyor belt operating in the solar convection zone. Further work with helioseismic inversions and a careful study of systematic effects are needed before firm conclusions of these large-scale flow structures can be made.
false
[ "meridional flow", "meridional flows", "mid-convection zone depths", "solar dynamics", "deep solar layers", "Global Oscillation Network Group Doppler data", "different measurement techniques", "meridional motions", "the solar convection zone", "decadal timescales", "Global Oscillation Network Group Doppler", "magnetic flux", "firm conclusions", "recent conclusions", "previous work", "partial evidence", "the convection zone", "a possible equatorward flow", "Travel-time differences", "travel-time differences" ]
11.317943
15.255902
2
678537
[ "Wang, Yi", "Xue, Wei" ]
2014JCAP...10..075W
[ "Inflation and alternatives with blue tensor spectra" ]
66
[ "Centre for Theoretical Cosmology, DAMTP, University of Cambridge, Cambridge CB3 0WA, U.K.; INFN-Sezione di Trieste, SISSA, via Bonomea 265, 34136 Trieste, Italy", "Centre for Theoretical Cosmology, DAMTP, University of Cambridge, Cambridge CB3 0WA, U.K.; INFN-Sezione di Trieste, SISSA, via Bonomea 265, 34136 Trieste, Italy" ]
[ "2014EPJC...74.3139C", "2014JCAP...07..020C", "2014JCAP...08..036M", "2014JCAP...09..012A", "2014JCAP...10..004S", "2014JCAP...10..029O", "2014JCAP...10..053B", "2014MPLA...2950185C", "2014NuPhB.885..196M", "2014PhLB..735..108C", "2014PhLB..736..488L", "2014PhLB..737...98A", "2014PhLB..738..206L", "2014PhRvD..90b3544H", "2014PhRvD..90d3505B", "2014PhRvD..90h3521L", "2014PhRvD..90j3515A", "2014PhRvD..90l3524H", "2014PhRvL.113c1301S", "2014SCPMA..57.1414C", "2014SCPMA..57.1431L", "2014SCPMA..57.1455L", "2014SCPMA..57.1466W", "2014arXiv1404.0238L", "2014arXiv1404.2222M", "2015ApJ...807..166S", "2015CQGra..32w4002B", "2015IJMPD..2441005D", "2015JCAP...02..003K", "2015JCAP...10..009G", "2015NuPhB.900..517C", "2015PhLB..745..105I", "2015PhRvD..92f3511C", "2015PhRvD..92l1304N", "2016JCAP...02..013C", "2016JCAP...07..030F", "2016JCAP...12..026B", "2016JCAP...12..040H", "2016NCimR..39..399G", "2016PhRvD..93f3005C", "2016PhRvD..93f3508C", "2016PhRvD..94d3002C", "2016Prama..86..353M", "2016arXiv160305836M", "2016arXiv160501615C", "2016arXiv161200998B", "2017IJMPD..2640004C", "2017JCAP...05..034D", "2017JHEP...04..058C", "2017PhRvD..95f3529M", "2017PhRvD..96f3506A", "2021Ap&SS.366..113M", "2021JCAP...06..010A", "2021PhRvD.103f3532B", "2022GReGr..54..156C", "2022PhRvD.105d3520B", "2023EPJC...83..816Z", "2023JHEAp..39...81V", "2023JHEP...07..146K", "2023PhRvD.108j3507B", "2024EPJC...84..278C", "2024JCAP...03..031B", "2024JCAP...06..004B", "2024JCAP...06..036D", "2024JHEP...02..008Y", "2024arXiv240412655P" ]
[ "astronomy", "physics" ]
29
[ "Astrophysics - Cosmology and Nongalactic Astrophysics", "General Relativity and Quantum Cosmology", "High Energy Physics - Phenomenology", "High Energy Physics - Theory" ]
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[ "10.1088/1475-7516/2014/10/075", "10.48550/arXiv.1403.5817" ]
1403
1403.5817_arXiv.txt
Inflation \cite{Guth81, Linde82} is the leading paradigm for the very early universe cosmology. Inflation has been proposed to explain the horizon, flatness and monopole problems in the standard hot big bang cosmology, and almost all the predictions of the simplest inflation model have now been tested. The observational tests of inflation includes \begin{itemize} \item Coherent and nearly scale invariant power spectrum of density perturbations. The power spectrum of the simplest slow roll inflation is \cite{Planck16} \begin{align} P_\zeta = \frac{H^2}{8\pi^2\epsilon M_p^2} \simeq 2.43 \times 10^{-9}~. \end{align} \item A small tilt of the scalar power spectrum. \begin{align} n_s - 1 = -2\epsilon - \eta \simeq 0.96~, \end{align} where $\eta \equiv \dot \epsilon / (H\epsilon)$ is the slow roll parameter defined from expansion. Now $n_s\geq 1$ is ruled out under the assumptions of simplest inflation models. \item Nearly Gaussian density fluctuations. The non-Gaussianities of the density fluctuations are tightly constrained at \begin{align} f_\mathrm{NL}^\mathrm{local} = 2.7 \pm 5.8~, \qquad f_\mathrm{NL}^\mathrm{equil} = -42\pm 75 \end{align} for the local shape and equilateral shape non-Gaussianities respectively. Those numbers indicate that, non-Gaussian components of the primordial fluctuations, even if exist, have to be at least 3$\sim$4 orders of magnitudes smaller than the Gaussian component. \item Gravitational waves. The recent BICEP2 experiment reports an over $5\sigma$ detection of gravitational waves \cite{Ade:2014xna}, with tensor to scalar ratio \footnote{There is a debate about whether the observed polarisation signal comes from primordial gravitational waves or dust contamination \cite{Mortonson:2014bja, Flauger:2014qra}.} \begin{align} r = 0.20^{+0.07}_{-0.05} ~(1\sigma \mathrm{CL})~. \end{align} This corresponds to a gravitational wave fluctuation amplitude \begin{align} P_\mathrm{T} = \frac{2H^2}{\pi^2 M_p^2} = 4.8 \times 10^{-10}~. \end{align} \end{itemize} Despite of the great success of inflation, there are still a few outstanding challenges for theorists and experimentalists. On the theoretical side, large field inflation is now favored. However, large field inflation is hard to construct from the effective field theory, and stringy UV completion points of view. The UV completion of inflation has long suffered from an $\eta$-problem \cite{Copeland:1994vg}, in which the mass of the inflaton is theoretically too large to allow enough e-folds of inflation. However, with the current data, a more serious $\epsilon$-problem emerges -- the observed energy scale of inflation is too high for an effective field theory or stringy model building to be under control. For single field slow roll inflation, at every e-fold, the inflaton rolls a distance of order $0.1 M_p$. In perturbative string theory, this field motion per e-fold is comparable with, or greater than the string scale $M_s$. As a result, one can no longer safely globally expand the inflaton field and ignore non-renormalizable terms. More discussions and a local reconstruction of the inflationary potential can be found in \cite{Ma:2014vua} On the observational side, there is yet another (and maybe the last in the foreseeable future, unless nature is so kind as to imprint other relics on the CMB sky or in the large scale structure) test for inflation which is possible in light of BICEP2, but not yet achieved -- the tilt of the tensor power spectrum. The simplest inflation models predict a consistency relation between $n_\mathrm{T}$ and $r$ as \begin{align} n_\mathrm{T} = - \frac{r}{8} = -0.025~. \end{align} Currently the data has not been good enough to test $n_\mathrm{T}$ precisely. However, there are a lot of ongoing and upcoming experiments in the near future \cite{Planck22, Ade:2014afa, Austermann12, Niemack10, Eimer12}, measuring $r$ at different scales, which provides a possibility for a precise measurement of $n_\mathrm{T}$. In this paper, we shall explore the possibility of blue $n_\mathrm{T}$. In Section \ref{sec:hints-blue-tensor}, the bound on $n_\mathrm{T}$ is derived the BICEP2 \cite{Ade:2014xna} and POLARBEAR \cite{Ade:2014afa} data. The string gas cosmology, null energy condition (NEC) violating inflation, and general initial condition are addressed in Sections \ref{sec:string-gas-cosmology}, \ref{sec:infl-viol-nec} and \ref{sec:infl-gener-init} respectively. In Section \ref{sec:other-possibilities}, a few other possibilities are discussed, including external sources for tensor modes, modified gravity and matter bounce. We conclude in Section \ref{sec:concl-disc}.
\label{sec:concl-disc} To conclude, we fit the tensor-to-scalar ratio and the tensor spectral tilt with data. The current data is not good enough to test the inflationary consistency relation but nevertheless blue tensor spectra are favored when all 9 bins of BICEP2 data, in combination with the POLARBEAR data, are used. From theoretical aspects, string gas cosmology predicts blue tensor spectra. However, the tilt is small, at the same order-of-magnitude of scalar tilt. Thus the future experiments targeting to test the inflationary consistency relation can also test this prediction of string gas cosmology. On the other hand, string gas cosmology predicts highly Gaussian density and tensor perturbations. This is unlike the other inflationary mechanisms, where a blue tilt also implies non-Gaussianities. The simplest model of G-inflation, on the other hand, tilts scalar and tensor power spectra in the same way. Thus with a red scalar spectral tilt, the tensor spectra are also red. In the parameter regime where both scale and tensor spectra are blue, the blueness of tensor perturbation is suppressed by slow roll parameter $\epsilon$. The equilateral non-Gaussianity of G-inflation, at $r=0.2$, is about $30$ and close to the current observational bound. The non-Gaussianities should be a model independent feature for super-inflation type models which generate blue tensor spectrum, because the perturbations have to behave differently from the background to avoid ghosts. Generalized initial conditions of inflation is left largely unconstrained. However, the generalized initial conditions are also sources of non-Gaussianities. Thus non-Gaussianities in the tensor sector would be a test of those class of models. Inflationary particle production is another possible source of tensor modes. We derived a bound of inflationary Hubble scale and energy density for this mechanism to work. Non-Gaussianities are also present in the case of particle productions. We hope to investigate the particle production mechanism and its relation between the blue tensor spectra in a future work. The current data shows hint of blue $n_\mathrm{T} \sim \mathcal{O}(1)$. The models with slow roll suppressed $n_\mathrm{T}$ would not be enough to explain such a hint. On the other hand, from modified initial conditions, external sources, inflation beyond slow roll, modified gravity and some models of matter bounce, blue and large $n_\mathrm{T}$ may be produced. It would be interesting to study those models in more details to see how the fitting of data is improved in those models. \noindent\textbf{Note added:} Two related works \cite{Gerbino:2014eqa, Ashoorioon:2014nta} appeared on arXiv on the same day as ours. \cite{Gerbino:2014eqa} (see also \cite{cosmocoffee}) overlaps with our Section \ref{sec:hints-blue-tensor}, and \cite{Ashoorioon:2014nta} overlaps with Section \ref{sec:infl-gener-init}. The data analysis of tensor tilt is also investigated by \cite{Cheng:2014bma}, \cite{Wu:2014qxa} and \cite{Cheng:2014ota}, within a few days before/after our paper. Among those papers, \cite{Cheng:2014bma} and \cite{Cheng:2014ota} reports a nearly zero central value, with small errorbars $\Delta n_\mathrm{T} \sim 0.48$ (BICEP2 only) and $\Delta n_\mathrm{T} \sim 0.24$ (BICEP2 + Planck + WP, with running of scalar spectral index). While the following works prefer blue tilt with considerably larger errorbars: \cite{Gerbino:2014eqa} (BICEP2 only), our result (BICEP2, BICEP2 + Planck + WP, POLARBEAR), \cite{cosmocoffee} (BICEP2 + Planck + WP and BICEP2 x Keck + Planck + WP) and \cite{Wu:2014qxa} (BICEP2 + Planck + WMAP + BAO).
14
3
1403.5817
We study the tilt of the primordial gravitational waves spectrum. A hint of blue tilt is shown from analyzing the BICEP2 and POLARBEAR data. Motivated by this, we explore the possibilities of blue tensor spectra from the very early universe cosmology models, including null energy condition violating inflation, inflation with general initial conditions, and string gas cosmology, etc. For the simplest G-inflation, blue tensor spectrum also implies blue scalar spectrum. In general, the inflation models with blue tensor spectra indicate large non-Gaussianities. On the other hand, string gas cosmology predicts blue tensor spectrum with highly Gaussian fluctuations. If further experiments do confirm the blue tensor spectrum, non-Gaussianity becomes a distinguishing test between inflation and alternatives.
false
[ "blue tensor spectrum", "blue tensor", "blue tensor spectra", "blue scalar spectrum", "blue tilt", "inflation", "string gas cosmology", "null energy condition", "general initial conditions", "the blue tensor spectrum", "alternatives", "Gaussian", "Gaussianity", "the inflation models", "Gaussianities", "highly Gaussian fluctuations", "the very early universe cosmology models", "BICEP2", "a distinguishing test", "the simplest G-inflation" ]
11.129851
-0.583047
89